{ "0405/astro-ph0405308_arXiv.txt": { "abstract": "{ We have detected longitudinal magnetic fields between 2 and 4\\,kG in three (WD\\,0446$-$790, WD\\,1105$-$048, WD\\,2359$-$434) out of a sample of 12 normal DA white dwarfs by using optical spectropolarimetry done with the VLT Antu 8 m telescope equipped with FORS1. With the exception of 40 Eri B (4\\,kG) these are the first positive detections of magnetic fields in white dwarfs below 30\\,kG. Although suspected, it was not clear whether a significant fraction of white dwarfs contain magnetic fields at this level. These fields may be explained as fossil relics from magnetic fields in the main-sequence progenitors considerably enhanced by magnetic flux conservation during the shrinkage of the core. A detection rate of 25 \\% (3/12) may indicate now for the first time that a substantial fraction of white dwarfs have a weak magnetic field. This result, if confirmed by future observations, would form a cornerstone for our understanding on the evolution of stellar magnetic fields. ", "introduction": "The major goal of this work is a better understanding of the role played by magnetic fields in the formation and evolution of stars. Magnetic fields are already an important ingredient during the collapse and fragmentation of protostellar clouds which ultimately determines the initial field of pre-main sequence stars. On the main sequence and at later stages of evolution the magnetic fields have a major impact on the angular momentum loss and stellar winds, on building-up chemical anomalies and abundance inhomogeneities across the stellar surface, on convection and the related coronal activity, and other evolutionary processes especially in interacting binaries. Very strong magnetic fields are detected in several white dwarfs and always present in pulsars \\citep{Mestel:01}. The first detection of a magnetic field on a white dwarf was made by \\cite{Kemp-etal:70} on Grw+$70^\\circ$ 8247, and large spectroscopic and polarimetric surveys have been carried out in the last two decades \\citep[\\eg\\ ][]{Hagen-etal:87, Reimers-etal:94, Schmidt-Smith:95, Putney:95, Kawka-etal:03}. For the white dwarfs the magnetic fields could simply be ``fossil'' remnants of the fields already present in main-sequence stars, but strongly amplified by contraction. This hypothesis assumes that the magnetic flux (\\eg\\ through the magnetic equator) is conserved to a large extent during the stellar evolution. According to the field amplification theory, the white dwarfs play an important role in the investigation of stellar magnetic fields. In main-sequence stars magnetic fields have been detected directly mainly for peculiar magnetic Ap and Bp stars with rather well organized fields and field strengths of the order $10^2-10^4$ Gauss. For weak fields in A to O stars ($B<10^2$ G) direct magnetic field detections are still very rare \\citep[\\eg\\ the field detections in early B stars reported by][] {Neiner-etal:03a,Neiner-etal:03b}. For sun-like stars ample evidence (coronal activity) for the presence of complicated small-scale fields exists, but direct measurements are only possible for the more active stars \\citep{Saar:96,Ruedi-etal:97,Valenti-Johns:01}. The contraction to a white dwarf amplifies the magnetic fields by about 4 orders of magnitude, so that weak and often undetectable magnetic fields on the main sequence become measurable during the white dwarf phase. This is supported by the known magnetic white dwarfs with megagauss fields ($B=10^6-10^9$~G). Their frequency and space distribution, as well as their mass, are consistent with the widely accepted view that they are the descendents of the magnetic Ap and Bp stars \\citep[\\eg\\ ][]{Mathys:01}. Another origin seems to be required for the magnetic degenerates with weaker fields (unless magnetic flux is lost during the contraction phase). Magnetic main-sequence stars with weaker magnetic fields have been suggested as their possible progenitor candidates \\citep{Schmidt-etal:03,Kawka-etal:03}. The B stars on which weaker fields have been detected may be the missing stars. However, even the most sensitive observations are limited to some tens of gauss on main-sequence stars. Thus, magnetic field amplification during stellar evolution may offer the opportunity to investigate $\\sim 1$~G magnetic fields (averaged global fields) in normal main-sequence stars with observations of $\\sim 1$~kG magnetic fields during the white dwarf stage. White dwarfs with magnetic fields below 100\\,kG have been either found by searching for circular polarization \\citep{Schmidt-Smith:94} or by looking for Zeeman splitting of narrow NLTE line cores in the Balmer lines, particularly in H$\\alpha$ \\citep{Koester-etal:98}. However, the splitting becomes undetectable in intensity spectra for weak fields ($<20$ kG) or for objects without narrow line core. Therefore, spectropolarimetry is the most promising technique for successful detections of weak magnetic fields. Up to now detections of magnetic fields below 30 kG have not been achieved, except for the very bright white dwarf 40 Eri B ($V=8.5$), in which \\cite{Fabrika-etal:03} have detected a magnetic field as low as 4\\,kG. The magnetic field detection limit can now be pushed down to a few kG for many white dwarfs with spectropolarimetry using 8-10~m class telescopes. In this paper we present and analyse VLT spectropolarimetric data of a sample of 12 white dwarfs in a search for weak magnetic fields. In Sect. 2 the observations and data reduction are described, while in Sect. 3 the observational method for obtaining the Stokes parameter ($V/I$) is described. Sect. 4 presents the method for determining weak magnetic fields analysing the circular polarisation due to a given magnetic field. In that section we also present the results of our analysis, along with the description of the $\\chi^2$-minimisation procedure applied to our data. The determination of the atmospheric and stellar parameters is presented in Sect. 5 and compared with those found in the literature. A discussion and conclusions are presented in Sect. 6. ", "conclusions": "In this work we have used the spectropolarimetric capability of the FORS1 instrument, together with the light collecting power of the VLT, in order to investigate the presence of magnetic fields in the range $1-10$~kG for a sample of 12 white dwarfs. Three of the stars out of 12 normal DA white dwarfs of our sample, WD\\,0446-789, WD\\,1105-048 and WD\\,2359-434, exhibit magnetic fields of a few kilogauss in one or all available observations (see Figs.\\,3, 4 and 5). The detection rate of 25~\\% suggests now strongly that a substantial fraction of white dwarfs have a weak magnetic field. With the exception of the bright white dwarf 40 Eri B, for which a magnetic field of only 4\\,kG had earlier been detected \\citep{Fabrika-etal:03}, WD\\,0446$-$789, WD\\,1105$-$048 and WD\\,2359$-$434 have the weakest magnetic fields detected so far in white dwarfs. In previous extended searches for weak fields in white dwarfs \\citep[\\eg][]{Schmidt-Smith:95} only very few objects have been found. Only six detections have been reported for magnetic fields below 100\\,kG, which are not all confirmed, and only three objects have a field weaker than about 50\\,kG. \\cite{Kawka-etal:03} have reported longitudinal magnetic fields in three stars but no significant detection was made because their $1\\sigma$ error was almost as large as the observed value itself. They concluded that the population of white dwarfs with magnetic fields in excess of 1\\,MG is well known, but that lower-field white dwarfs remained undetected. Note, however, that our investigation is based on the averaged longitudinal component of the magnetic field, meaning that the maximum magnetic field at the white dwarf surface can be stronger, depending on the field geometry (described \\eg\\ by offset dipoles, or more complex distributions; with the underestimate being larger for a more complex magnetic distribution) and on the orientation relative to the observer. Therefore, our results for the three objects with a positive detection are lower limits, since cancellation effects are expected. The population of known magnetic white dwarfs presently comprises some 125 stars \\citep{Wickramasinghe-Ferrario:00,Gaensicke-etal:02,Schmidt-etal:03}. \\cite{Wickramasinghe-Ferrario:00} and \\cite{Jordan:01} established that about 3\\% of all white dwarfs had a magnetic field above 100\\,kG. Results from the first two years of the Sloan Digital Sky Survey \\citep{Gaensicke-etal:02,Schmidt-etal:03} show the total number of known magnetic white dwarfs with $B \\geq 1.5$ MG being 6\\%. Recently, \\cite{Liebert-etal:03} have found that the incidence of magnetism at the level of $\\sim$ 2\\,MG or greater is at least $\\sim$10\\%, or higher. They suggest that the total fraction of magnetic WDs may be substantially higher than 10\\% due to the limited spectropolarimetric analyses capable of detecting lower field strengths down to $\\sim$ 10\\,kG. Our 3 detections out of 12 objects seem to indicate that low magnetic fields on white dwarfs ($<$10\\,kG) are frequent while high magnetic fields are relatively rare. However, with only three detections this hypothesis remains insecure. If confirmed by future observations, the investigation of weak magnetic fields in white dwarfs could form a cornerstone for the future investigation of the properties and evolution of stellar magnetic fields. Our sample of white dwarfs is too small to discuss in detail the dependence of the magnetic field strength on the stellar parameters (masses and cooling ages). It is, however, worthwhile to mention that two of our detections (WD\\,0446$-$789 and WD\\,1105$-$048) have masses of only 0.5 \\Msolar. This means that their progenitors on the main-sequence had less than 1 \\Msolar\\ \\citep{Weidemann:2000}. These two stars are therefore very different from the majority of white dwarfs with megagauss magnetic fields which tend to have higher masses \\citep{Greenstein-Oke:82,Liebert:88} and, therefore, high-mass parent stars. Measurements of weak magnetic fields are now possible for many white dwarfs with the new large telescopes, which allow a magnetic field function (MFF, in analogy to the mass function) to be constructed in the $1-100$~kG range once a sufficient number of detections have been made. Such a MFF can be compared to the corresponding function for main-sequence stars \\citep{Bychkov-etal:97} and will provide input for answers to the following key questions on the evolution of magnetic fields in stars: Are the magnetic fields in white dwarfs just the fossil relics of magnetic main-sequence stars strengthened by contraction due to conservation (to a large extent) of magnetic flux? Or do the magnetic fields develop considerably through the final stages of stellar evolution? Are the strongly magnetic white dwarfs a distinctive class of objects or do they just represent a tail of the distribution of magnetic fields present in all white dwarfs? Is there a dependence between magnetic field strength and mass as found in the case of magnetic WDs with higher field strengths? Do the magnetic field strengths correlate with temperature, which would be a hint for a decay on the white dwarf cooling sequence? Several authors have suggested that the frequency of magnetic white dwarfs may increase with decreasing effective temperature, luminosity and with increasing cooling age \\citep[\\eg\\ ][]{Valyavin-Fabrika:98,Liebert-etal:03}, and may decrease sharply with distance \\citep[][]{Fabrika-Valyavin:98}. Alternatively to the fossil origin of the magnetic field in white dwarfs, \\cite{Markiel-etal:94} and \\cite{Thomas-etal:95} have shown that a weak magnetic field of $\\approx 1.3$\\,kG in the variable DB star GD\\,358 can be explained by an $\\alpha\\omega$ dynamo. The magnetic field in this star has been inferred indirectly by analyzing the $g$-mode oscillation spectrum taken with the WET \\citep[Whole Earth Telescope,][]{Winget-etal:94}. However, according to the atmospheric parameters, the convection zone in all of our sample stars should be too shallow to support an $\\alpha\\omega$ dynamo. Another long-standing problem in white dwarf research is the issue why metals are accreted by helium-rich white dwarfs in the range $8000\\,\\kelvin < T_{\\rm eff} < 15000$\\,\\kelvin\\ during the passage through an interstellar cloud while almost no hydrogen is brought into the white dwarf atmosphere. \\cite{Illarionov-Sunyaev:75} suggested that fields below $10^5$\\,G provide a screening mechanism to separate hydrogen and ionized species from grains in white dwarfs accreting from the interstellar matter. Since this phenomenon always occurs in this type of star, it is possible that at this level all white dwarfs contain magnetic fields. \\cite{Friedrich:04} have searched for circular polarisation in one DBZ and one DBAZ which have accreted metals, but three or four orders of magnitude less hydrogen than expected. In one case (L745-46A) a magnetic field of 7\\,kG (1 $\\sigma$ error of $\\pm 2$\\,kG; 99\\%\\ confidence interval of $\\pm 6$\\,kG) was found, which, however, was based on the H$\\alpha$ line only. For the second object (GD\\,40), only an upper limit of 12\\,kG (99\\% confidence) could be derived from polarization measurements around three spectral lines. Theoretically, magnetic fields of 4\\,kG and 250\\,kG, respectively, would be required for the screening by the propeller mechanism to be efficient around the two stars. At our signal-to-noise ratio, magnetic fields down to about 2 \\,kG can be measured. There is still the possibility that all magnetic white dwarfs contain surface magnetic fields at the 1\\,kG level. In order to test this hypothesis, much longer exposure times would be necessary, even with the VLT." }, "0405/astro-ph0405622_arXiv.txt": { "abstract": "Most old distant radio galaxies should be extended X--ray sources due to inverse Compton scattering of Cosmic Microwave Background (CMB) photons. Such sources can be an important component in X-ray surveys for high redshift clusters, due to the increase with redshift of both the CMB energy density and the radio source number density. We estimate a lower limit to the space density of such sources and show that inverse Compton scattered emission may dominate above redshifts of one and X-ray luminosities of $10^{44}$\\ergps, with a space density of radio galaxies $> 10^{-8}$ Mpc$^{-3}$. The X-ray sources may last longer than the radio emission and so need not be associated with what is seen to be a currently active radio galaxy. ", "introduction": "{\\em Chandra} has revealed extended X-ray emission from a wide range of radio sources out to high redshifts. Jets and the lobes and cocoons of radio quasars and galaxies have been imaged with unprecedented resolution (e.g. Chartas et al 2000, Schwarz et al 2001, Harris \\& Krawczynski 2002, Kataoka et al 2003, Comastri et al 2003, Wilson, Young \\& Shopbell 2001, Kraft et al 2002). At low redshifts, extended emission directly associated with the radio lobes is seen through its inverse Compton emission in some objects (e.g. Fornax A; Feigelson et al 1995, Kaneda et al 1995; 3C219, Comastri et al 2003). In the powerful radio source Cygnus A, diffuse X-ray emission is also detected from the radio cocoon, i.e. the reservoirs of shocked material associated with the radio expansion (Wilson, Young \\& Smith 2003). When however a modest radio source lies in a rich cluster, the surface brightness of the inverse Compton emission in soft X-rays can be so low that it cannot be separated from thermal emission and the lobes appear as holes in the X-ray emission due to displacement of the hot gas by the lobes (e.g. the Perseus cluster, Fabian et al 2000, Sanders et al 2004; Hydra A, McNamara et al 2000; A2052, Blanton et al 2001). Extended X-ray emission is also associated with an increasing number of radio sources at cosmological distances, for example 3C~294 (Fabian et al 2003), 3C~9 (Fabian, Celotti \\& Johnstone 2003), PKS\\,1138-262 (Carilli et al 2002), 4C~41.17 (Scharf et al 2003), GB 1508+5714 (Yuan et al 2003, Siemiginowska et al 2003), although often the low photon count rate makes it hard to disentangle the different X-ray components. In general however, much of the emission can be interpreted as due to inverse Compton scattering of non-thermal radio emitting electrons on CMB photons (Felten \\& Rees 1967, Cooke, Lawrence \\& Perola 1978, Harris \\& Grindlay 1979). Where directly associated with powerful jets relativistic bulk motion may be involved (Celotti, Ghisellini \\& Chiaberge 2001; Tavecchio et al 2000). The steep increase in the energy density of the CMB with redshift $z$ (as $(1+z)^4$) partially compensates for the large distance to such sources (Felten \\& Rees 1967, Schwartz 2002), thereby making them detectable. Independently of the spatial distribution, the presence of extended radio synchrotron emission is a direct indication of the presence of a non-thermal population of relativistic particles. These particles, at least, must produce high energy emission via inverse Compton scattering of CMB photons (direct measurements and upper limits for such emission are used to estimate the intracluster magnetic field). Here, we consider the effective number density of extended X-ray emitting sources due to this process as a function of X-ray luminosity and show that they can be a serious contaminant to X--ray surveys searching for clusters and protoclusters at high redshift. The outline of the paper is as follows: in Section 2 we estimate the ratio of (synchrotron) radio to (inverse Compton) X-ray emission, while in Section 3 we estimate the corresponding X-ray luminosity functions of radio sources as a function of $z$ and compare them with those of X-ray clusters. A discussion in Section 4 concludes this Letter. A cosmology with $\\Omega_{\\Lambda}=0.7, \\Omega_{\\rm M}=0.3$, $H_0=50$ km s$^{-1}$ Mpc$^{-1}$ has been assumed. ", "conclusions": "The above estimates indicate that powerful radio galaxies are expected to be found in significant numbers at redshift $z\\approxgt 1$ as extended X-ray sources. Indeed, we expect that most old distant radio galaxies are also extended X--ray sources. It should be stressed that for any $B< B_{\\rm x}$ the above estimates provide a {\\em lower} limit on both the individual luminosity and the number density of extended X--ray emitting radio sources for the following reasons. Firstly, extended X-ray emission can also be produced by non-thermal particles which do not contribute to the 151 MHz emission, further increasing the actual X-ray luminosity with respect to the estimates given above: lobes, cocoons, relics and jets can emit not only as inverse Compton emission on the CMB, but also via other emission processes such as synchrotron self-Compton, bremsstrahlung and inverse Compton on other photon fields, such as far-infrared photons in the vicinity of the massive, extremely luminous galaxies detected in the sub-mm at high redshifts. Secondly, the radio-emitting electrons cool faster than the X-ray emitting ones (for $B< B_{\\rm x}$). Thus the radio luminosity function might significantly underestimate, by a factor corresponding to the relative cooling times $t_{\\rm cool} (\\gamma_{\\rm R})/t_{\\rm cool} (\\gamma_{\\rm x})$, the number density of sources and the volume pervaded by non-thermal electrons with energy $\\sim \\gamma_{\\rm x}$ contributing to the X-ray emission. In other words the estimates above refer to 'prompt' X-ray emission only over the cooling timescale due to radio emission. Taking into account this ratio effectively increases the normalization of the luminosity function of radio sources by factors $\\sim 3-10$ (see Fig. 1b). In this respect it is interesting to notice that indeed in the case of 3C294 (Fabian et al 2001) the X-ray emission extends much further than the radio structure (visible at 5 GHz), indicating that $B 3\\ 10^{-16}$ erg cm$^{-2}$ s$^{-1}$. A potential diagnostic of nonthermal emission is emission at energies $\\gg1$ keV is also expected. Although a study of the effective extension of the particle distribution requires a knowledge of the acceleration processes at work as well as its history (see Sarazin 1999), if a source is radio emitting at 151 MHz, for any given $B$ the spectrum produced via the scattering of CMB photons extends at least up to $\\sim 40 (1+z) B^{-1}_{-6}$ keV; possibly to a few hundreds MeV (for $\\gamma \\sim 10^6$). Because of photon redshifting, detection of any emission above a few keV at high $z$ would be a clear signature of a non-thermal component, presumably due to inverse Compton emission. Interestingly, one out of the six extended X-ray sources detected by Bauer et al. (2002) has a high temperature with respect to the $L_{\\rm x}-T$ cluster correlation. The inverse Compton emission is a direct result, and also a probe, of the major non-gravitational energy injection phase of the present day intracluster medium. In fact, the radio luminosity of radio galaxies grossly underestimates the intrinsic power of their jets which can be 1000 or more times greater. The radio-emitting plasma is likely to be confined by some intracluster or intragroup gas, which is displaced outward. The energy dumped into the immediate surroundings of such sources can be considerable and thereby influence the gaseous properties of clusters and groups (Ensslin et al 1997, Valageas \\& Silk 1999; Wu, Fabian \\& Nulsen 2000). Integrating the radio galaxy luminosity function over time leads to a comoving energy input of about $10^{57}$\\ergpMpc (Inoue \\& Sasaki 2001) which can (pre)heat -- although the precise mechanism is not yet clear -- the intracluster medium by 1--2~keV per particle, so explaining much of the non-gravitational scaling behavior of groups and clusters (e.g. Lloyd-Davies et al 2000). The intracluster or intragroup medium will be highly disturbed during the energy injection phase and for up to a core crossing time (about a Gyr) after. This means that the detection and interpretation of Sunyaev-Zeldovich signals from high redshift clusters (Carlstron, Holder \\& Reese 2003) may be complicated. Alternatively, the lack of detection of a large population of non-thermal extended X-ray emitters would provide interesting information about the radio source/cluster magnetic field evolution. This would suggest that the radio emitting particles have lower energy than the radio emitting ones, indicating $B>B_{\\rm x}$, and thus point to positive evolution in the magnetic field associated with the diffuse radio emission at higher redshifts, although it would be difficult to precisely quantify and interpret such result. \\noindent" }, "0405/astro-ph0405078_arXiv.txt": { "abstract": "We present a spectroscopic sample of 48 early-type galaxies in the rich cluster Abell~2218 and 50 galaxies in Abell~2390. Since both samples are very similar, we combine them and investigate a total number of 98 early-type galaxies at $z\\sim 0.2$. A subsample of 34 galaxies with \\textit{HST} structural properties is used to construct the Fundamental Plane. Elliptical and S0 galaxies show a zeropoint offset of $\\overline{m}_{r}\\sim0.43$~mag with respect to the local Coma FP. Both sub-samples, ellipticals and lenticulars, exhibit a similar, mild evolution and small scatter. The moderate amount of luminosity evolution is consistent with stellar population models of passive evolution, if $z_{f}\\ge2$ is assumed. ", "introduction": "Clusters of galaxies are powerful laboratories in investigating the formation and evolution of galaxies. Rich clusters are dominated by early-type (E+S0) galaxies. A detailed study of the properties of these galaxies provides beneficial insights in their formation and evolution. In particular, the paradigm of hierarchical galaxy formation can be critically tested. In the local Universe numerous investigations revealed that E+S0 form a very homogeneous galaxy population (e.g., \\cite{BLE92}; \\cite{BBF93}). Their structural parameters (effective surface brightness $\\langle\\mu\\rangle_e$ and the size as described by the effective radius $R_e$) and kinematics (velocity dispersion $\\sigma$) represent a tight correlation in three dimensional parameter space, the Fundamental Plane (\\cite{DD87}). Furthermore, they show a small scatter in their relations of colours (e.g., Mg--($B-V$)), $M/L$ ratios and absorption line indices with velocity dispersion (e.g., Mg--$\\sigma$). However, the question arises whether E+S0s are truly one single family or rather a diverse group with different formation and evolutionary processes. Therefore, one aim of this work is to explore if there are differences between the properties of elliptical and S0 galaxies. Previous spectroscopic studies were limited to a small number of the more luminous galaxies. To overcome bias and selection problems of small samples, we focus in this study of the clusters Abell~2218 (\\cite{Z2001}) and Abell~2390 on a large number of objects ($N=98$), spanning a wide range in luminosity, in case of A\\,2390 $21.4$ 0.0). Such sources have been identified by \\citet{kobulnicky99c} and \\citet{johnson03b} as extremely young, dense, heavily-embedded star clusters (or ``ultradense \\HII\\ regions''). Since the sample of galaxies searched for such regions remains small, the total fraction of starbursting systems that contain such extreme star formation regions is not well known. However, preliminary estimates suggest these sources may enshroud $\\sim$ 10\\% of the O-star populations of starburst galaxies \\citep{johnson01}. If radio continuum emission is observed at multiple frequencies and at matched resolutions, we can discern which of the above physical processes dominates within individual star formation regions. By concentrating on nearby starburst galaxies that have resolved stellar and cluster populations in the visual or infrared, we can probe the youngest phases of star formation, from heavily-embedded star clusters to the supernova remnants resulting from massive star evolution. This provides detailed insights into the processes that regulate star formation in starburst galaxies, including feedback and cluster formation and evolution timescales. These are important quantities both locally and in the high-redshift universe. Here, we present new VLA observations of the dwarf starburst galaxy NGC\\,625. This system is a nearby \\citep[D $=$ 3.89\\,$\\pm$\\,0.22 Mpc;][]{cannon03} prototypical dwarf starburst that has revealed many striking properties in recent multiwavelength investigations (see Table~\\ref{t1} for a summary of basic galaxy parameters). NGC\\,625 currently hosts a massive starburst with a star formation rate $\\sim$ 0.05 \\msun\\,yr$^{-1}$ and which displays Wolf-Rayet (W-R) emission features ({Skillman, C{\\^ o}t{\\' e}, \\& Miller 2003a,b}\\nocite{skillman03b,skillman03a}). We have performed a recent star formation history analysis using spatially resolved stellar HST/WFPC2 photometry \\citep{cannon03}, and find that the current burst is actually long-lived (\\gsim\\ 50 Myr) compared to canonical expectations based on W-R star populations (ages \\lsim\\ 6 Myr; {Conti 1991}\\nocite{conti91}; {Schaerer, Contini, \\& Kunth 1999a}\\nocite{schaerer99a}). This extended burst of star formation appears to have disrupted the \\HI\\ disk, and this system is currently undergoing outflow of \\HI\\ from the major starburst region {(Cannon \\etal\\ 2004a)}\\nocite{cannon04a}. This outflow is also seen in diffuse x-rays \\citep{bomans98} and in \\ion{O}{6} absorption from FUSE spectroscopy (Cannon \\etal\\ 2004b, in preparation). \\placetable{t1} From the above properties it is clear that the star formation in NGC\\,625 is violent and is having a dramatic impact on the ISM and potentially on the surrounding IGM as well. Since this galaxy is one of only a small subset of star-forming dwarfs that demonstrate such extreme properties, it is important that we explore other observational avenues with which to better understand its evolution. Here we use multifrequency VLA radio continuum data to infer the dominant emission processes in the major star formation regions, as well as various characteristics of the recent star formation in this galaxy. ", "conclusions": "\\label{S5} We have presented new VLA radio continuum imaging, in the BnA and CnB arrays at L, C and X-bands, of the nearby dwarf starburst galaxy NGC\\,625. The global spectral index is nearly flat, suggesting that thermal emission dominates the radio continuum luminosity. Examination of the spectral indices of individual star formation regions shows a mix of thermal and nonthermal processes. The highest-surface brightness \\HII\\ regions NGC\\,625\\,A and NGC\\,625\\,B show strong free-free emission, while the less-luminous \\HII\\ region NGC\\,625\\,C is dominated by nonthermal synchrotron radiation. At $\\sim$ 2\\arcsec\\ ($\\sim$ 40 pc) resolution, we do not find any deeply-embedded sources in NGC\\,625. However, we do find (with only 3\\,$\\sigma$ significance) one low-luminosity source that has no obvious optical counterpart and that is located in the region of highest optical extinction. We interpret the mix of thermal and nonthermal emission from the main star formation regions as an age progression, with free-free emission arising from \\HII\\ regions with ages $<$ 10 Myr, and synchrotron radiation prominent in regions older than this. Comparing to a limited sample of well-studied dwarf starburst and W-R galaxies, we find that global thermal emission is common in these types of systems. The spectral indices are typically flatter than in \\HII\\ galaxies, and this is again suggestive that the dominant free-free component in dwarf starburst and W-R galaxies implies younger starburst ages than in more evolved systems where synchrotron emission dominates. The importance of small-scale processes (i.e., feedback) in regulating the evolution of emission in these systems is highlighted. While a simple age interpretation is suitable for most of this small nearby sample, more complex interpretations are needed when comparing multiwavelength properties of various galaxies. The formation and evolution of outflows is an important parameter, but is relatively poorly constrained by observations. More detailed modeling and higher-resolution observations of larger samples of dwarf starburst and W-R galaxies would be beneficial for understanding the evolution of the radio continuum properties of these important star-forming systems." }, "0405/astro-ph0405528_arXiv.txt": { "abstract": "During the transition from a neutral to a fully reionized universe, scattering of cosmic microwave background (CMB) photons via free-electrons leads to a new anisotropy contribution to the temperature distribution. If the reionization process is inhomogeneous and patchy, the era of reionization is also visible via brightness temperature fluctuations in the redshifted 21 cm line emission from neutral Hydrogen. Since regions containing electrons and neutral Hydrogen are expected to trace the same underlying density field, the two are (anti) correlated and this is expected to be reflected in the anisotropy maps via a correlation between arcminute-scale CMB temperature and the 21 cm background. In terms of the angular cross-power spectrum, unfortunately, this correlation is insignificant due to a geometric cancellation associated with second order CMB anisotropies. The same cross-correlation between ionized and neutral regions, however, can be studied using a bispectrum involving large scale velocity field of ionized regions from the Doppler effect, arcminute scale CMB anisotropies during reionization, and the 21 cm background. While the geometric cancellation is partly avoided, the signal-to-noise ratio related to this bispectrum is reduced due to the large cosmic variance related to velocity fluctuations traced by the Doppler effect. Unless the velocity field during reionization can be independently established, it is unlikely that the correlation information related to the relative distribution of ionized electrons and regions containing neutral Hydrogen can be obtained with a combined study involving CMB and 21 cm fluctuations. ", "introduction": "The large angle bump in the polarization-temperature cross correlation power spectrum \\cite{Zaldarriaga:1996ke} measured in WMAP data \\cite{Bennett:2003bz} indicates an optical depth to electron scattering of 0.17 $\\pm$ 0.04 \\cite{Kogetal03}. To explain both this high optical depth and the Lyman-$\\alpha$ optical depth from Gunn-Peterson troughs \\cite{GunPet65} towards $z \\sim 6$ quasars in the Sloan Digital Sky Survey \\cite{Fan:2001vx} requires a complex reionization history \\cite{Cen:2003ey,Chen:2003sw,Cen:2002zc}. While slight modifications to the large angular scale polarization power spectra exist with different reionization histories that integrate to the same electron scattering optical depth, one cannot use these changes to fully reconstruct the reionization history as a function of redshift due to large cosmic variance associated with anisotropy measurements at a few tens degree angular scales \\cite{Kaplinghat:2002vt}. Under standard expectations for reionization, mainly due to UV light emitted by first luminous objects, the reionization process is expected to be both patchy and inhomogeneous \\cite{Barkana:2000fd}. This leads to fluctuations in the electron scattering optical depth and to a modulation of both temperature and polarization contributions such that new anisotropy fluctuations are generated at arcminute scales corresponding to inhomogeneities in the visibility function \\cite{Hu:1999vq,Santos:2003jb}. The increase in sensitivity and angular resolution of upcoming CMB anisotropy data suggests that such small-scale fluctuations can soon be used to understand the reionization history and associated physics beyond measurements related to large-scale polarization alone. The secondary reionization-related anisotropies, unfortunately, are expected to be confused with other secondary effects such as the Sunyaev-Zel'dovich (SZ; \\cite{SunZel80}) contribution at late times associated with galaxy clusters and first supernovae \\cite{Oh:2003sa}, and higher order effects such as gravitational lensing \\cite{Hu:ee}. Though this confusion can be partly removed, such as through frequency cleaning in the case of SZ \\cite{Cooray:2000xh} or subtraction of lensing via higher order statistics \\cite{Okamoto:2003zw}, secondary anisotropies alone cannot be used to extract the detailed history of reionization beyond the integrated optical depth \\cite{Santos:2003jb,Zhang:2003nr}. An interesting possibility involves the cross-correlation between small-scale CMB polarization maps and images of the high redshift universe. This cross-correlation can help distinguish broad aspects such as whether the universe reionized once or twice \\cite{Cooray:2003dt}. Beyond the reionization history, mainly in the scattering visibility function as a function of redshift, it would be interesting to study additional details related to the reionization process such as the relative distribution of free-electrons and neutral Hydrogen. Since the brightness temperature fluctuations associated with the 21 cm spin-flip transition \\cite{Fie58} trace the neutral Hydrogen distribution, the combination involving secondary CMB and 21 cm fluctuations \\cite{Tozzi:1999zh} could then provide additional details related to reionization. In this paper, we consider such a combined study in the form of a cross-correlation between arcminute scale CMB anisotropy maps and images of the high redshift universe around the era of reionization from the 21 cm background. We make the assumption that the reionization process is inhomogeneous and patchy such that at certain epochs during partial ionization, contributions are generated to both the 21 cm background and CMB. In such an era, we also assume that the spin temperature of the neutral Hydrogen distribution is decoupled from CMB and is dominated by the kinetic temperature; In this case, the signature in 21 cm emission is fluctuations related to the distribution of neutral Hydrogen, though, in the limit where the spin temperature is smaller than that of CMB (prior to reionization and appearance of first sources), 21 cm fluctuations will be tightly coupled to that of large angular scale CMB temperature; In this case, one naturally expects a perfect cross-correlation between these two maps. We do not consider this possibility as the reionization process is expected to heat the IGM by $z \\sim 20$. Also, opportunities for very low frequency observations, where 21 cm signatures from redshifts prior to reionization are expected, are extremely limited. In the era of partial reionization, the existence of the proposed cross-correlation is due to the fact that inhomogeneities that lead to fluctuations in both CMB (in terms of the electron distribution) and 21 cm background (via the neutral content) trace the same underlying density field. One expects the two to spatially correlate though this would be an anti-correlation as regions containing free-electrons will be mostly free of neutral material. This fact is captured by a spatial cross-power spectrum between free-electron and neutral Hydrogen fluctuations. Since the 21 cm background allows one to probe the power spectrum of neutral Hydrogen alone, while CMB probes the power spectrum of reionized patches, in combination, the cross-power spectrum provides additional information on physics related to reionization. Unfortunately, while there is a strong (anti) cross-correlation between free-electrons and neutral Hydrogen fluctuations, we find the observable projected angular cross-spectrum to be insignificant due to a geometric cancellation. This cancellation comes from the CMB side and involves the line of sight projection of the velocity field during reionization. The same cross-correlation can also be studied using a higher order correlation in the form of a bispectrum in Fourier space involving the large scale velocity field of ionized regions from the Doppler effect, arcminute scale CMB anisotropies, and the 21 cm background. This measurement avoids the geometric cancellation associated with the line of sight projection, but due to the large cosmic variance associated with the velocity field traced by the Doppler effect, its measurement is limited to signal-to-noise ratios of order ten (with noise dominated maps) or, at most, a hundred. We discuss the measurement of proposed cross-correlations using CMB maps from upcoming missions, such as Planck surveyor and the South Pole Telescope (SPT), and maps of the 21cm background with, say, the Square Kilometer Array. Unfortunately, with signal-to-noise ratios around ten or below, it is unlikely that one can use the proposed cross-correlation bispectrum to easily extract detailed information on the relative distribution between neutral Hydrogen and electrons. The paper is organized as follows. In \\S~\\ref{sec:deriv}, we derive the existence of the cross-correlation both in terms of the cross power spectrum between CMB temperature and 21 cm fluctuations and a higher order bispectrum that avoids a geometric cancellation associated with the cross power spectrum. In \\S~\\ref{sec:results}, we discuss our results and suggest that though there is adequate signal-to-noise to perform a cross correlation study in terms of the bispectrum. We briefly discuss how this measurement can be improved and what information related to reionization can be extracted from this measurement. We conclude with a summary in \\S~\\ref{sec:summary}. ", "conclusions": "\\label{sec:summary} During the transition from a neutral to a fully reionized universe, scattering of cosmic microwave background (CMB) photons via free-electrons lead to a new anisotropy contribution to the temperature distribution. If the reionization process is inhomogeneous and patchy, the era of reionization is also visible via brightness temperature fluctuations in the redshifted 21 cm line emission of neutral Hydrogen. Since regions containing electrons and neutral Hydrogen are expected to trace the same underlying density field, the two are (anti) correlated and this is expected to be reflected in the anisotropy maps in terms of a cross-correlation between arcminute-scale CMB temperature and the 21 cm background. In terms of the angular cross-power spectrum, unfortunately, this correlation is insignificant due to a geometric cancellation associated with second order CMB anisotropies. Thus, it is unlikely that the cross-correlation spectrum between small-scale CMB and 21 cm fluctuations will be measurable even with maps involving a high signal-to-noise per pixel. The same cross-correlation between ionized and neutral regions, however, can be studied using a bispectrum involving large scale velocity field of ionized regions from the Doppler effect, arcminute scale CMB anisotropies involving reionization signal, and the 21 cm background. While the geometric cancellation is partly avoided, the signal-to-noise ratio related to this bispectrum is reduced via large cosmic variance related to the velocity fluctuations traced by the Doppler effect. Unless velocity fluctuations can be independently established, it is unlikely that the correlation information related to the relative distribution of ionized electrons and regions containing neutral Hydrogen can be obtained with a combined study involving CMB and 21 cm fluctuations." }, "0405/astro-ph0405144_arXiv.txt": { "abstract": "We discuss the very different methods in each wavelength band for selecting and finding Active Galactic Nuclei (AGN). We briefly review the history of the different techniques for finding AGN and compare and contrast the advantages and difficulties of selection in different wavelength bands. We stress the strong selection effects in each wavelength band and the difficulty of defining complete samples. Of all the techniques presently used, we conclude that selection in the hard X-ray band via imaging and spectroscopy is the most complete and allows the best estimate of the number and evolution of active galaxies. However, all of the techniques have difficulties at low luminosities where emission due to stellar processes can have similar sizes and luminosities. ", "introduction": "Looking at the 60 year history of observations of active galaxies, it is clear that the definition of what they are has strongly influenced the methods of finding them. From our present perspective, many of the techniques used over the past 40 years are not truly appropriate and are more in the line of the famous joke of the drunk looking under the lamp post for his lost car keys. In this chapter, I will use the words Active Galactic Nuclei (alias AGN or quasars) to be the equivalent of radiating supermassive black holes, even though this perspective is very recent. The difficulty in finding AGN is defining what makes the observed\\footnote{While very frequently the inferred emitted radiation is rather different from that intrinsically produced by the region around the black hole, the nature of surveys is such that we must rely on the observed properties of these sources in order to find and identify them.} radiation different from that due to other processes, in particular those related to normal stars and stellar evolution (e.g., supernovae).\\footnote{To date, all the surveys for AGN have relied on detection of radiation across the electromagnetic spectrum. Perhaps in the distant future we will be able to search for AGN via neutrinos, gravitational waves, or even very high-energy cosmic rays, but this is still quite uncertain.} This has often been a process of exclusion; that is, the emission does not resemble that from stars or stellar processes. Dust, high-luminosity emission from starbursts, and the possible effects of unusual types of stars complicate the issue. The strong effects of observing in different spectral ranges also need to be taken into account; for example, at $R=22$ there are only 100 ``optically-selected'' AGN per square degree, but at the equivalent flux level in the $2-10$~keV X-ray band of $3\\times 10^{-15}$~ergs~cm$^{-2}$~s$^{-1}$, there are 1000~deg$^{-2}$. Finally, it is clear that the ``non-stellar'' signature has a wide variety of forms that gives rise to the ``zoo'' of names for active galaxies. The spectral energy distributions, optical emission-line properties (strengths, widths, and nature), line of sight column densities, time variability characteristics, and bolometric luminosities of Seyfert~1 galaxies, Seyfert~2 galaxies, BL Lacertae objects, LINERs (Low-Ionization Nuclear Emission Regions), and quasars (to use the names of the largest samples of objects) are all rather different (see, e.g., Risaliti \\& Elvis, this volume). It has taken many years and a large amount of effort to finally come to the realization that all these classes are manifestations of the same underlying physical process---emission from near to a supermassive black hole. However, even today it is not certain if all of these sources are driven solely by accretion, or whether there is also energy extraction from the spin of the black hole (see Armitage, this volume). It is also not clear if the energy production is dominated by radiation, relativistic particle production, or bulk motion of material. It is entirely possible that there is a simple relation between the ``names'' of the sources and their physical natures, but at present this seems very complex and not unique. As opposed to stellar classifications, there is not a one-to-one relationship between the class of the object and its physical nature. However, there are some clear distinctions: for example, in BL Lacertae objects, the observed radiation is dominated by emission from relativistic particles in a jet in our line of sight, and in Seyfert~2 galaxies, the line of sight to the central source is blocked by large amounts of dust and gas (see Hewett \\& Foltz 1994 for an earlier review and a detailed discussion of the many systematic effects in quasar surveys). Before discussing the field in general, it is important to consider what a survey really is. As Hewett \\& Foltz (1994) point out, there are three types of surveys: (1) those that find objects, (2) those that find objects consistently, and (3) those with well-defined selection criteria that allow probabilities to be assigned for selection as a function of survey parameters. Surveys of the first type are the easiest, since the goal is only to provide sources for study that meet some criteria. Surveys of the second type are homogenous in their properties, but completeness is not important. Many of the issues discussed below are more or less important depending on which type of survey is being performed. It is only surveys of the third type that allow comparisons to be made of different wavelength regimes and different survey techniques. While these surveys are the most scientifically important, they are also the most difficult to do. \\subsection{A Short History of AGN Search Techniques} From a historical perspective (e.g., Osterbrock 1991), the first indications of ``non-stellar'' activity in the nuclei of galaxies came from the discovery of strong, broad emission lines in NGC1068 (Fath 1913; Slipher 1917) and the discovery of the jet in M87 (Curtis 1917). The spectral features found in NGC1068 are almost never found in stars or supernovae remnants and thus were an indication of some new phenomena occurring. It took 50 years before the morphology of the ``jet'' in M87 was related to non-stellar processes.\\footnote{We now know that jets can also occur in stellar processes (for example, in Herbig-Haro objects, young neutron stars [e.g., see the {\\em Chandra\\/} image of the Vela pulsar], supernovae remnants [e.g., see the {\\em Chandra\\/} image of Cassiopeia A], and the X-ray and radio emission from some luminous galactic X-ray sources [e.g., Cygnus X-3]). It seems as if the jet phenomenon is related to the emission of large amounts of energy over a short period of time and into a restricted volume.} The first ``sample'' of non-stellar activity was that of Seyfert in 1943, who found a wide variety of strong ``broad'' lines in the nuclei---but not elsewhere---of several otherwise ``normal'' galaxies. It was clear from this early paper that there was something quite unusual about these sources and that they must be fairly common, but it took almost 20 more years for significant progress to be made. This next step occurred with the discovery of extragalactic radio sources and attempts to find optical counterparts for them. The discovery of low to moderate redshift ``active galaxies'' as the optical counterparts to several bright radio sources (Baade \\& Minkowski 1954) showed a new type of ``active galaxy''. It was clear even in 1958 (Minkowski) that there was tremendous scatter in the optical properties of the identifications at a fixed radio flux. Schmidt (1963), Greenstein \\& Matthews (1963), and Schmidt \\& Matthews (1964) discovered that the optical counterparts of several luminous radio sources were stellar-looking sources at large redshifts, and thus were very luminous, compact, extragalactic sources with non-stellar spectra. They were subsequently named ``quasars'' for quasi-stellar radio sources. The optical properties of these radio sources were very similar to each other, indicating that a class of objects had been found. Sandage (1965) realized that there were sources with the same general optical properties as quasars that were not radio sources. These sources had blue colors, meaning a large ultraviolet (UV) flux relative to the classical optical band, fairly high variability in their continuum intensity, and most had strong, broad emission lines over a wide range of ionization. It was rapidly realized that the nuclei of some ``Seyfert galaxies'' had similar properties to quasars (Woltjer 1959; Burbidge, Burbidge, \\& Sandage 1963), and for the last 30 years, these sources have been grouped under the name Active Galactic Nuclei or AGN (I think that the first use of this name in the literature is from Burbidge 1970). However, even early on, it was clear that not all AGN resembled quasars. There were Seyfert~2 galaxies, which do not have broad lines or strong non-stellar continua but do have strong, narrow forbidden lines that could not be produced by ionization from normal stars. Also, there were BL Lacertae objects, which usually do not have optical or UV emission lines but do have a very strong non-thermal continuum and a wide variety of optical colors. It was thus clear from the start that identifying complete samples of AGN would require a wide variety of techniques and criteria. Outside of trying to identify the optical counterparts to the radio sources (it took over 30 years to completely identify all of the sources in the 3CR survey, the first large radio survey; Spinrad et al.\\ 1985), the first systematic search for AGN that I can find in the literature was the realization (Arp et al.\\ 1968) that the Markarian survey of compact galaxies with blue stellar colors contained a large number of sources with the properties of Seyfert galaxies. At almost the same time, Sargent (1970) used similar critieria for sources from the Zwicky survey and also found numerous AGN. These two early works, combined with the radio surveys, set the standard for AGN surveys: using photographic techniques to find sources with compact, blue nuclei and following up with optical spectroscopy (see Weedman 1977 for an early review). However, these searches were completely empirical; i.e., they were looking for sources that had the properties of sources that one already knew were active galaxies. It is clear that AGN have a very wide range of relative parameters, from the line strengths and line widths (ranging down to sources without any emission lines at all; e.g., BL Lacertae objects), to the continuum colors, to the amplitudes and timescales of variability. Thus, having an inclusive definition is very difficult. It is fair to say that most workers have had a difficult time coming up with an AGN definition that is totally complete and not subject to noise and strong selection effects. Recent large optical surveys, such as the Two Degree Field (2dF) and the Sloan Digital Sky Survey (SDSS), have focused on well-defined color or line strength criteria that allow them to be well defined but clearly incomplete. ", "conclusions": "" }, "0405/astro-ph0405372_arXiv.txt": { "abstract": "The current standard theory of the origin of the Moon is that the Earth was hit by a giant impactor the size of Mars causing ejection of iron poor impactor mantle debris that coalesced to form the Moon. But where did this Mars-sized impactor come from? Isotopic evidence suggests that it came from 1AU radius in the solar nebula and computer simulations are consistent with it approaching Earth on a zero-energy parabolic trajectory. But how could such a large object form in the disk of planetesimals at 1AU without colliding with the Earth early-on before having a chance to grow large or before its or the Earth's iron core had formed? We propose that the giant impactor could have formed in a stable orbit among debris at the Earth's Lagrange point $L_4$ (or $L_5$). We show such a configuration is stable, even for a Mars-sized impactor. It could grow gradually by accretion at $L_4$ (or $L_5$), but eventually gravitational interactions with other growing planetesimals could kick it out into a chaotic creeping orbit which we show would likely cause it to hit the Earth on a zero-energy parabolic trajectory. This paper argues that this scenario is possible and should be further studied. ", "introduction": "\\label{Section:1} The currently favored theory for the formation of the Moon is the giant impactor theory formulated by Hartmann $\\&$ Davis (1975) and Cameron $\\&$ Ward (1976). Computer simulations show that a Mars-sized giant impactor could have hit the Earth on a zero-energy parabolic trajectory, ejecting impactor mantle debris that coalesced to form the Moon. Further studies of this theory include (Benz, Slattery, $\\&$ Cameron 1986, 1987; Benz, Cameron, $\\&$ Melosh 1989; Cameron $\\&$ Benz 1991; Canup $\\&$ Asphaug 2001; Cameron 2001, Canup 2004; Stevenson 1987). We summarize evidence favoring this theory: (1) It explains the lack of a large iron core in the Moon. By the late time that the impact had taken place, the iron in the Earth and the giant impactor had already sunk into their cores. So, when the Mars-sized giant impactor hit the Earth in a glancing blow, it expelled debris, poor in iron, primarily from mantle of the giant impactor which eventually coalesced to form the Moon (cf. Canup 2004, Canup2004B). Computer simulations (assuming a zero-energy parabolic trajectory for the impactor) show that iron in the core of the giant impactor melts and ends up deposited in the Earth's core. (2) It explains the low (3.3 grams/cm$^3$) density of the Moon relative to the Earth (5.5 grams/cm$^3$), again due to the lack of iron in the Moon. (3) It explains why the Earth and the Moon have the same oxygen isotope abundance - the Earth and the giant impactor came from the same radius in the solar nebula. Meteorites originating from the parent bodies of Mars and Vesta, from different neighborhoods in the solar nebula have different oxygen isotope abundances. The impactor theory is able to explain the otherwise paradoxical similarity between the oxygen isotope abundance in the Earth combined with the difference in iron. This is perhaps its most persuasive point. (4) It explains, because it is due to a somewhat unusual event, why most planets (like Venus and Mars, Jupiter and Saturn) are singletons, without a large moon like the Earth. Competing ideas have not had comparable success. For example, the idea that the Earth and the Moon formed together as sister planets in the same neighborhood fails because it doesn't explain the difference in iron. Whereas the idea that the Moon formed elsewhere in the solar nebula and was captured into an orbit around the Earth fails because its oxygen isotope abundances would have to be different. That a rapidly spinning Earth could have spun off the Moon(from mantle material) is not supported by energy and angular momentum considerations, it is argued. Still, the giant attractor theory has some puzzling aspects. Planets are supposed to grow from planetesimals by accretion. How did an object so large as Mars, form in the solar nebula at exactly the same radial distance from the Sun without having collided with the Earth earlier, before it could have grown so large. Indeed, such must have been the case during the formation of Venus and Mars for example. It's also hard to imagine an object as large as Mars forming in an eccentric Earth-crossing orbit. One might expect large objects forming in the solar nebula to naturally have nearly circular orbits in the ecliptic plane, like the Earth and Venus. Besides, a Mars-sized object in an eccentric orbit would not be expected to have identical oxygen abundances relative to the Earth, and would collide with the Earth on a hyperbolic trajectory not the parabolic trajectory that the successful computer simulations of the great impact theory have been using.(Recent collision simulations by Canup(2004) place an upper limit of 4 km/s for the impactor's velocity-at-infinity approaching the Earth, setting an upper limit on its eccentricity of $\\stackrel{<}{\\sim} 0.13$.) The Mars-sized object needs to form in a circular orbit of radius 1 AU in the solar nebula but curiously must have avoided collision with the Earth for long enough for its iron to have settled into its core. Is there such a place to form this Mars-sized object? Yes: the Earth's Lagrange point $L_4$(or$L_5$) which is at a radius of 1 AU from the Sun, with a circular orbit $60^o$ behind the Earth(or $60^o$ ahead of the Earth for $L_5$). After the epoch of gaseous dissipation in the inner solar nebula has passed we are left with a thin disk of planetesimals interacting under gravity. The three-body problem shows us that the Lagrange point $L_4$(or equivalently $L_5$) for the Earth is stable for a body of negligible mass even though it is maximum in the effective potential. Thus, planetesimals can be trapped near $L_4$ and as they are perturbed they will move in orbits that can remain near this location. This remains true as the Earth grows by accretion of small planetesimals. Therefore, over time, it might not be surprising to see a giant impactor growing up at $L_4$(or $L_5$). In the Discussion section we argue that there are difficulties in having the giant impactor come from a location different from $L_4$(or $L_5$). Examples of planetsimals remaining at Lagrange points of other bodies include the well known Trojan asteroids at Jupiter's $L_4$ and $L_5$ points. As another example, asteroid 5261 Eureka has been discovered at Mars' $L_5$ point. (There are five additional asteroids also thought to be Mars Trojans: 1998 $VF_{31}$, 1999 $UJ_7$, 2001 $DH_{47}$, 2001 $FG_{24}$, 2001 $FR_{127}$.) The Saturn system has several examples of bodies existing at the equilateral Lagrange points of several moons, which we discuss further in the \"Note added\" after the Discussion section. We propose that the Mars-sized giant impactor can form as part of debris at Earth's $L_4$ Lagrange point. (It could equally well form at $L_5$, but as the situation is symmetric, we will simply refer in the rest of the paper, unless otherwise indicated, to the object forming at $L_4$; the argument being the same in both cases). As the object forms and gains mass at $L_4$, we can demonstrate that its orbit about the Sun remains stable. Thus, it has a stable orbit about the Sun, and remaining at $L_4$ keeps it from collision with the Earth as it grows. Furthermore, this orbit is at exactly the same radius in the solar nebula as the Earth so that its oxygen isotope abundances should be identical. It is allowed to gradually grow and there is time for its iron to settle into its core, and the same also happens with the Earth. The configuration is stable providing the mass of the Earth and the mass of the giant impactor are both below .0385 of the mass of the Sun, which is the case. But eventually, we numerically demonstrate that gravitational perturbations from other growing planetesimals can kick the giant impactor into a horseshoe orbit and finally into an orbit which is chaotically unstable in nature allowing escape from $L_4$. The giant impactor can then enter an orbit about the Sun which is at an approximate radial distance of 1 AU, which will gradually creep toward the Earth; leading, with large probability, to a nearly zero-energy parabolic collision with the Earth. Once it has entered the chaotically unstable region about $L_4$, a collision with the Earth is likely. We will discuss this phenomenon in detail in Section 4. For references on the formation of planetesimals and related issues, see (Goldreich 1973; Goldreich $\\&$ Tremaine 1980; Ida $\\&$ Makino 1993; Rafikov 2003; Wetherill 1989). We are considering instability of motion near $L_4$ due to encounters by planetesimals. (The instability of the Jupiter's outer Trojan asteroids due to the gravitational effects of Jupiter over time studied in Levinson, Shoemaker E. M. $\\&$ Shoemaker C. S. (1997), is a different process.) Horseshoe orbits connected with the Earth exist. In fact, an asteroid with a 0.1 km diameter, 2002 $AA_{29}$, has recently been discovered in just such a horseshoe-type orbit which currently approaches the Earth to within a distance of only 3.6 million km (Conners et al. 2002). Horseshoe orbits about the Sun of this type are also called Earth co-orbiting trajectories, which are in 1:1 mean motion resonance. An interesting pair of objects in horseshoe orbits about Saturn are discussed in the Note after the Discussion section. A theoretical study of the distribution of objects in co-orbital motion is given by Morais $\\&$ Morbidelli (2002). In this paper we describe a special set of collision orbits with the Earth which exist due to escape from $L_4$ due to planetesimal perturbations. The perturbations cause a gradual peculiar velocity increase of the mass forming at $L_4$ so that it eventually achieves a critical escape velocity to send it toward a parabolic Earth collision approximately in the plane of motion of the Earth about the Sun. The region in velocity space where escape from $L_4$ occurs in this fashion is relatively narrow. This mechanism therefore involves a special set of $L_4$ ejection trajectories which creep towards collision with the Earth. At the end of Section 3 we will present a full simulation in three-dimensions of a collision of a Mars-sized impactor with the Earth assuming a thin planetesimal disk, using the general three-body problem, where planetesimal encounters with both the impactor and the Earth are done in a random fashion. The Appendix of this paper discusses the dynamics of the random planetesimal encounters. The paper has several main results: \\noindent We show that a stable orbit at $L_4$ exists where a Mars-sized giant impactor could grow by accretion without colliding with the Earth. We show that eventually perturbations by other planetesimals can cause the giant impactor to escape from $L_4$ and send it onto a horseshoe orbit and then onto a creeping chaotic trajectory with an appreciable probability of having a near parabolic collision with the Earth. In the Discussion section we argue how this scenario fits in extremely well with giant impactor theory and explains the identical oxygen isotope abundances of the Earth and the Moon. The solar system itself provides a testing ground for our model. As we have mentioned the Trojan asteroids show that planetesimals can remain trapped at Lagrange points, and in the \"Note added\" we point out that the system of Saturn's moons provide examples where the phenomenon we are discussing can be observed, supporting our model. Finally both in the \"Note added\" and in the Appendix we discuss prospects for future work. The spirit of this paper is to suggest the intriguing possibility that the hypothesized Mars-sized impactor could have originated at $L_4$(or$L_5$). It is hoped that this lays the ground work for more detailed simulations and work in the future. ", "conclusions": "\\label{ref:Conclusions} \\medskip \\medskip We have shown that the giant impactor could have formed at $L_4$($L_5$) and then escaped on a creeping chaotic trajectory to impact the Earth, with a near parabolic encounter in agreement with simulations. We note that there are difficulties if the giant impactor came from a location other than $L_4$(or $L_5$). To illustrate this assume that it came from elsewhere. Since the Earth's orbit and Venus' are nearly circular and co-planar even at the current epoch, after 4.5 billion years of perturbations, this suggests that the early disk of planetesimals in the neighborhood of the Earth was quite thin and that the planetesimals in the disk were in orbits that had low eccentricity (e) and inclination (i), $e \\sim i \\ll 1$. The critical impact parameter for collision with the Earth for a small planetesimal is $$b_m = r_E(1+[V_{es}^2/V_{pec}^2])^{1/2},$$ where $V_{pec}$ is the peculiar velocity of the planetesimal, i.e., $V_{pec}\\sim e \\times V_{orb} \\sim e \\times 30$ km s$^{-1}$, and $V_{es}$ is the escape velocity from the surface of the Earth (see Appendix). If $V_{pec}/V_{orb} < 0.004$, then $b_m \\sim r_E(V_{es}/V_{pec}) > e \\times 1AU \\sim i \\times 1AU \\sim (V_{pec}/V_{orb})1AU$, and the planetesimals whose semi-major axes are within a distance of $b_m$ of the Earth's distance from the Sun of 1AU will likely suffer collision with the Earth within a short number of years since we expect the orbits to be chaotic, and the impact parameter with the Earth is less than the critical impact parameter $b_m$ for collision with the Earth. This will clear out a region of $\\pm b_m$ around 1AU {\\it Except for planetesimals in stable orbits around $L_4$(or$L_5$)}. Planetesimals at nearly 1AU from the Sun--and not at $L_4$(or$L_5$)--will quickly be accreted by the Earth, before having the chance to grow large by accretion themselves. T.R. Cowley (lecture at Univ. of Michigan 2002) has noted this problem, saying that \"Advocates of the Big Whack hypothesis usually say that the impactor must have been formed near the Earth. This is neither probable nor impossible. It is not probable because the Earth could have readily swept up materials that would have formed the other body. It is not impossible because we do not know the precise conditions of the accumulation of the Earth, and cannot say how improbable assembly of the putative impactor near one astronomical unit really was.\") The stable location at $L_4$(or $L_5$) answers the probability question, offering a reasonably likely scenario for forming the giant impactor near 1AU without the material first being swept up by the Earth. Once this cleared out region of $\\pm b_m$ has been established, there will be no further quick accretion onto the Earth, because the planetesimal's orbits will not take them to within an impact distance $b_m$ from the Earth. Then they will have to diffuse in by two-body relaxation-from perturbations by other planetesimals and planets. This two-body relaxation process will slowly put planetesimals into the gap region again and there will be be quick accretion from the gap. The giant impactor is expected to be one of the later impactors to hit the Earth because the successful simulations of the formation of the Moon start with the Earth already at nearly its current mass, showing that its subsequent accretion (after the giant impactor hit) is assumed to be small (Canup $\\&$ Asphaug 2001). The giant impactor should also be expected to be one of the latter impacts because planetesimals, including the Earth, grow by accretion with time and that would have also allowed more time for the giant impactor to have grown by accretion itself. If the giant impactor is one of the {\\it latter} impactors as argued by Canup $\\&$ Asphaug (2001) (after most accretion for the Earth has been completed) then if it is not from $L_4$(or $L_5$) it must originally come from {\\it either} significantly outside 1AU or significantly inside 1AU. But then it would violate one of the key advantages of the great impactor theory: namely, point (3) in the Introduction, which explains why the Earth and Moon have the same oxygen isotope abundance- namely that the Earth and the giant impactor came from the same radius in the solar nebula. Meteorites from different neighborhoods in the solar nebula(those associated with parent bodies of Mars and Vesta for example) have different oxygen isotope abundances. The impactor theory is able to explain the otherwise paradoxical similarity between the oxygen isotope abundance in the Earth and the Moon combined with the difference in iron. The Earth has oxygen isotope abundances that are an average over all the planetesimals it has accreted--some initially from inside 1AU and some from outside. A giant impactor forming outside 1AU and drawn in by two-body interactions would have oxygen isotope abundances intermediate between Earth and Mars and therefore not identical with the Earth. Standard giant impact theory has the Moon formed primarily out of mantle material from the giant impactor material (see (Canup 2004)). The rest of the material in the giant impactor is absorbed by the Earth and the iron core of the giant impactor eventually finds its way into the Earth's core, leaving the Moon iron depleted relative to the Earth. The Moon has been found to have a small core and this is assumed to be from giant impactor material. (See (W\\\"anke 1999) for a discussion of how the giant impactor theory can accommodate this.) If the Moon derives from giant impactor material then it would have, the theory proposes, isotopic abundances identical with the Earth if it was formed near 1AU and this is observed to be the case (Clayton $\\&$ Mayeda 1996; Wiechert et al 2001; W\\\"anke 1999; Lodders $\\&$ Fegley 1997 . But if the giant impactor came from significantly outside outside 1AU its isotopic abundances would be significantly different from that of proto-Earth. Furthermore, since $M_I$ is only 10$\\%$ the mass of the Earth, this would pollute the proto-Earth's isotopic abundances with only a 10$\\%$ contribution from the giant impactor. This would give the Earth and the Moon different isotopic abundances, if the giant impactor came from significantly outside 1AU. A similar trouble occurs if the giant impactor originated significantly inside 1AU, if the oxygen isotope abundances inside 1AU are heterogeneous as well. (At present we have no meteorites in our possession whose parent bodies are thought to be Mercury or Venus. So we currently have no data for oxygen abundances inside 1AU.) On the other hand, consider what happens if the giant impactor originated at $L_4$(or $L_5$). It is in a stable orbit, so it is not immediately accreted onto the Earth, and can grow large and hit the Earth later, alleviating the problem mentioned by Cowley. It sits nicely at 1AU and accretes exactly the same type of material the Earth does, some diffusing from outside 1AU, and some from inside. The integral of the oxygen isotope abundances of the accretion should be identical with that of the Earth. Eventually, perturbations kick the giant impactor out of its stable orbit and it collides quickly with the Earth. When the giant impactor hits the Earth and kicks out the Moon, since the Earth and giant impactor have identical isotope ratios, the Earth and Moon should have identical isotope abundances even though the Earth and Moon are polluted to different extents by giant impactor material. This is an advantage to the giant impactor model, producing automatic agreement with proposition (3) of the giant impactor model. Since this is one of the latter accretion events for the Earth in terms of the accumulation of its mass, the oxygen isotope abundances for the Earth and Moon will not be further significantly changed by post-giant impactor accretion. Thus we propose the following scenario. \\medskip Debris remains at $L_4$ (as the Trojan asteroids prove). From this debris a giant impactor starts to grow like the Earth through accretion as described above. As the forming giant impactor reaches a sufficient mass ($\\sim .1 m_{Earth}$), it gradually moves away from $L_4$ through gravitational encounters with other remaining planetesimals and it randomly walks in peculiar velocity. It gradually moves farther and farther from $L_4$ approximately on the Earth's orbit in a horseshoe orbit at 1AU, until it acquires a peculiar velocity of approximately 180 m/s. The giant impactor then performs breakout motion where it performs a number of cycles about the Sun, repeatably passing near to the Earth. In a time span roughly on the order of 100 years it collides with the Earth on a near parabolic orbit. \\medskip \\medskip We present here a mechanism for the origin of a Mars-sized Earth impactor and describe the path it would take to arrive at Earth collision via a special class of slowly moving chaotic collision trajectories. The analysis shows that Earth collision along these trajectories is likely. Approaches for further work are discussed in the Appendix. \\medskip \\medskip \\medskip \\noindent {\\it Note added in proof} \\medskip As we have discussed, the giant impactor could have grown up in a stable orbit at Earth's $L_4$ (or $L_5$) point where a stable orbit is possible and an object could remain and be able to grow by accretion without hitting the Earth early-on. We expect this phenomenon could occur when there was a thin disk of planetesimals (in nearly circular orbits). We have noted that Saturn's rings are an example of such a thin disk of planetesimals(in this case, chunks of ice plus some dirt) observable today. Saturn's regular icy moons (inside the orbit of Titan) are all in nearly circular orbits of low eccentricity suggesting that they formed out of a thin disk of planetsimals(ice chunks) rather like Saturn's rings today only larger in extent. In such a situation we might expect our scenario to operate. Therefore it is quite interesting that we can find examples of objects at $L_4$ (or $L_5$), or escaping from $L_4$ (or $L_5$) in the Saturn system. Saturn's moon Helene co-orbits at the $L_5$ point (60$^o$ ahead) of the larger moon Dione. Helene has a largest diameter of 36 km and Dione has a diameter of 1120 km. Saturn's moons Telesto (diameter 34 km) and Calypso (diameter 34 km) occupy both the $L_4$ and $L_5$ points relative to Saturn's moon Tethys (diameter 1060 km). We would say that Helene, Telesto and Calypso originated in a planetesimal disk(of ice chunks) at these stable Lagrange points and have grown in place there surviving till the present without colliding with Dione or Tethys. The rest of the planetesimals(ice chunks) have accreted onto the regular moons of Saturn.(Saturn's rings themselves lie inside the Roche limit where the formation of large objects is forbidden by accretion.) While these Lagrange moons are small relative to the primary, growth of larger objects with respect to the primary is also possible. Saturn's moons Epimetheus (119 km diameter) and Janus (179 km diameter) co-orbit in horseshoe orbits just like the one we found for the giant impactor near breakout (Figure 5). We would say that Epimetheus formed at a Lagrange point of Janus and grew along with it by accretion from the planetesimal disk. Later perturbations by other planetesimals kicked it out into a horseshoe orbit just short of breakout. Thus, an object (Epimetheus) nearly as large as the primary (in this case Janus) can form and end up in a horseshoe orbit. Just a little more perturbation and Epimetheus would achieve breakout and likely collide with Janus. These provide examples of the phenomena described in this paper that can be observed today. A similar pair of co-orbiting objects in horseshoe orbits in another solar system could be easily detected using stellar radial velocity data. This would appear to be a planet in circular orbit about the star whose mass was observed to mysteriously vary. The mass variation would be approximately sinusoidally in time with a period significantly longer than the orbital period of the primary. For example, if the secondary had a mass 0.1 times that of the primary (like the giant impactor) then this would show up as a nearly sinusoidal variation of 10$\\%$ in the deduced mass of the primary. If the two were nearly equal in mass there would be a 100$\\%$ variation in the mass. (When they were near each other at one end of the horseshoe orbit the effective mass perturbing the star would be nearly doubled, and when they circulated to be on opposite sides of the star their perturbation would temporarily vanish.) We should have a look among the known cases of extra-solar planets for such cases. Granted, we are currently able to see only gas giant planets(which may have even migrated inward) rather the terrestrial ones we are considering, but still it would be interesting to look. If one found such a case, it would be easy to prove. Also of particular interest is the Earth co-orbiting asteroid 2002 $AA_{29}$ which is in a horseshoe orbit relative to the Earth. Of course, in addition to the giant impactor there can be other Lagrange-point debris particularly at the other stable Lagrange point not occupied by the giant impactor. This material may have been kicked out early-on by other planetesimals or by the giant impactor itself as it escaped into a horseshoe orbit. Does any of this material survive to the present day? 2002 $AA_{29}$ (diameter $<$ 0.1 km) is in a horseshoe orbit at 1AU, virtually identical to the horseshoe orbits found by us in Figure 5. This asteroid approaches the Earth closely (3.6 million miles away) once every 95 years while circling the sun at 1AU. It was near one of these close approaches that it was discovered in 2002. After a number of cycles, it is briefly captured for a period of 50 years as a quasi-satellite of the Earth, before returning to the 95-year horseshoe orbit cycle. Its rather large inclination (10.7$^o$) saves it from collision with the Earth. This object may have originated near $L_4$(or $L_5$), and have been kicked out into a horseshoe orbit (perhaps by the giant impactor itself). If that is so, it could be composed of the same material that also formed the seeds for the Earth and the giant impactor. A sample return from this asteroid thus offers the possibility of obtaining some primordial material from the same reservoir that produced the Earth and the Moon. In this case, it should have oxygen isotope abundances similar to those found for the Earth and the Moon and an iron abundance similar to that of the Earth. The final oxygen isotope abundances and iron abundances of the Earth reflect not only their seed material (originally from 1AU) but also the integral of the abundances accreted later, from material originally inside and outside 1AU. Thus, any slight differences in oxygen isotope abundances would be helpful in illuminating the accretion process. It would of course be very interesting to measure the age of a 2002 $AA_{29}$ sample. On the other hand, if the sample has oxygen isotope abundances identical with the Earth and the Moon, but is poor in iron like the Moon, that would suggest it was part of the splash material kicked out by the giant impactor at near escape velocity which did not coalesce onto the Moon but rather ended circling the Sun at 1AU and then became trapped at $L_4$(or $L_5$) where it moved for perhaps a considerable time before finally being kicked out. If the sample has completely different oxygen isotope abundances from those of the Earth and Moon, that would indicate an origin elsewhere in the solar nebula (not at 1AU) and we would then have to explain how it somehow got perturbed into a low eccentricity horseshoe orbit at 1AU. (Most Earth-crossing asteroids perturbed into their current orbits from the main belt should have much larger eccentricities according to Ipatov and Mather 2002.) Bottke, et al (1996) have previously suggested that low eccentricity objects near 1AU could have an origin tracing back to the Earth-Moon system, and radar results suggest (Ostro, et al 2003) that 2002 $AA_{29}$ has a high albedo which supports this hypothesis (according to Connors et al 2004). A sample return from asteroid 2002 $AA_{29}$ is thus of particular scientific interest and may provide important clues as to the origin of the Earth and the giant impactor that formed the Moon. \\medskip \\noindent {\\bf Acknowledgments} \\medskip \\noindent We would like to thank Scott Tremaine and Peter Goldreich for helpful comments, and also Robert Vanderbei for use of his solar system simulator. \\medskip \\noindent Partial support for this work for J. Richard Gott, III is from NSF grant AST-0406713, and for Edward Belbruno from grants by NASA, Office of Space Science, and Goddard Space Flight Center. \\newpage" }, "0405/hep-ph0405186_arXiv.txt": { "abstract": "The heavy gravitino in the minimal supergravity (mSUGRA) models is likely to be the lightest supersymmetric particle (LSP). Produced from the late decays of the metastable Weakly Interacting Massive Particles (WIMPs) such as the lightest neutralinos, the stable gravitinos can be plausible candidates for the cold dark matter in the universe. Such gravitino dark matter can naturally evade the current detection experiments due to its superweak couplings. However, this scenario must be subjected to the constraints from the Big Bang nucleosynthesis (BBN) predictions for light element abundances as well as the Wilkinson Microwave Anisotropy Probe (WMAP) data for the relic density. Assuming the popular case in which the lightest neutralino is the next-to-lightest supersymmetric particle (NLSP), we find that requiring BBN predictions for light element abundances to agree with the WMAP data can impose upper and lower mass bounds on both the gravitino LSP and the neutralino NLSP. A scan over the mSUGRA parameter space, subjected to the BBN constraints, the WMAP data and the $b\\to s \\gamma$ bounds, shows that the low $\\tan\\beta$ ($\\lsim$ 40) region as well as the region accessible at CERN Large Hadron Collider (LHC) will be severely constrained. Such stringent constraints on the parameter space might be instructive for testing this scenario in future collider experiments. ", "introduction": "The nature of the dark matter is one of the mysteries in today's physical science. It has been intensively explored both theoretically and experimentally. Studies showed that the cosmic dark matter is plausibly composed of non-baryonic Weakly Interacting Massive Particles (WIMPs) \\cite{review}. While the Standard Model of particle physics cannot provide a candidate for the dark matter WIMP, the popular supersymmetric theory with R-parity conservation can provide a good candidate, i.e., the lightest supersymmetric particle (LSP). So far the widely studied scenario is that the lightest neutralino is assumed to be the LSP. However, despite of the overwhelming popularity of this scenario, other possibilities should not be ignored due to the following reasons. On the one hand, the success of such neutralino dark matter scenario may be spoiled by the problems caused by gravitino in the reheating era \\cite{spoil}. On the other hand, the neutralino dark matter scenario has not yet been confirmed by current experiments\\cite{sudan}. One possible scenario other than neutralino LSP is that the gravitino is assumed to be the LSP. Such gravitino LSPs can form warm or cold cosmic dark matter, depending on the gravitino mass: \\begin{itemize} \\item[(i)] In some low-energy SUSY breaking models, like the gauge mediated SUSY breaking (GMSB) models, the gravitino can be as light as KeV, much lighter than other supersymmetric particles. It can thus form the warm dark matter. Note that the recent WMAP data imposed severe constraints on the dark matter type. As analyzed in \\cite{warm}, while a very tiny component of dark matter can be the hot neutrinos, the warm dark matter is ruled out due to the detected early re-ionization of the universe at a redshift $z \\approx 0.20$. Therefore, the scenario of warm dark matter gravitino is not favored by recent observation. \\item[(ii)] In the popular mSUGRA models, the gravitino mass is unspecified and only known to be of the weak-scale. Such heavy gravitino is possibly the LSP and can form the cold dark matter in the universe. In contrast to the highly constrained scenarios, in which the gravitino is produced as a thermal relic \\cite{relic} or produced during reheating \\cite{reheat}, a new scenario, assuming the gravitino to be produced from the late decays of the thermal relic WIMPs, was recently proposed in \\cite{feng}. \\end{itemize} In this article we focus our attention on this new gravitino dark matter scenario. Since the heavy gravitinos couple gravitionally, they are naturally the so-called Superweakly Interacting Massive Particles (SuperWIMPs). As a plausible candidate for the cold dark matter in the universe \\cite{feng}, the gravitino SuperWIMP can naturally evade the current dark matter detection experiments due to its superweak couplings. However, this scenario must be subjected to the constraints from the Big Bang nucleosynthesis (BBN) as well as the WMAP data \\cite{wmap}: \\begin{itemize} \\item The late decays of WIMPs (like neutralinos) into gravitino SuperWIMPs will release electromagnetic (EM) and hadronic energy. Such energy release will alter the BBN predictions for light element abundances \\cite{ellis,hadronic}. Requiring the resulted predictions for light element abundances to agree with the measured values will impose strong constraints on the gravitino dark matter scenario. \\item WMAP precisely measured many quantities, especially the total matter density and the baryon density, \\begin{eqnarray} \\Omega_m h^2 = 0.135_{-0.009}^{+0.008} \\ , ~~~ \\Omega_b h^2 = 0.0224_{-0.009}^{+0.009} \\ . \\end{eqnarray} From such results we can deduce the 2$\\sigma$ range for the cold dark matter density \\begin{eqnarray} \\label{density} \\Omega_{CDM} h^2 = 0.1126_{-0.0181}^{+0.0161}\\ , \\end{eqnarray} which is dramatically more accurate than previous results and agrees quite well with other approaches. Such precise measurements will impose strong constraints on the parameter space of gravitino dark matter scenario. \\end{itemize} The aim of this article is to examine these constraints on the parameter space of mSUGRA in this new gravitino dark matter scenario. Assuming the popular case that the lightest neutralino is the NLSP, we will examine the constraints on the mSUGRA parameter space from the BBN light element abundances, the WMAP data of relic density as well as the $b\\to s \\gamma$ branching ratio data. ", "conclusions": "We examined the constraints on the newly proposed dark matter scenario, in which the gravitino is assumed to be the LSP and produced from the late decays of metastable NLSP. Although such gravitino dark matter can naturally evade the current detection experiments due to its superweak couplings, we found, however, that this scenario is subjected to stringent constraints from the BBN predictions for light element abundances as well as the WMAP data for the relic density. Assuming the popular case that the lightest neutralino is the NLSP in mSUGRA models, we found that the low $\\tan\\beta$ ($\\lsim$ 40) region as well as the region accessible at the LHC are severely constrained. The popular mSUGRA models will be explored in future colliders like the LHC. For this purpose, it is important to know which part of the parameter space is viable and thus should be primarily explored. In this regard, the stringent cosmological constraints on the mSUGRA parameter space obtained in this work will be useful. Especially, if the ongoing and planned dark matter detection experiments fail to find any dark matter signal, it will imply that the dark matter interactions are too weak and thus the gravitino dark matter scenario will be favored. Then, to test this scenario at colliders, the stringent cosmological constraints on the parameter space will be quite instructive. This would serve as a good example that the studies in cosmology and astrophysics can shed some light on particle collider physics. On the other hand, the LHC could explore mSUGRA parameter space up to $m_{1/2}\\sim 1400$ GeV (700 GeV) for small (large) values of $m_0$, assuming 100 $fb^{-1}$ of integrated luminosity \\cite{LHC}. If the LHC results finally restrain the parameter space to one of the regions obtained in this work, then it implies that the gravitino dark matter scenario is favored. In this sence, the studies in particle physics can provide some insights in the understanding of dark matter in cosmology. We address that our study in this work is just illustrative instead of exhaustive. We assumed the popular case that the lightest neutralino is the NLSP in mSUGRA models. Actually, other super particles, like tau-slepton, are also likely to be the NLSP in mSUGRA models. If tau-slepton is assumed to be the NLSP, there are some theoretical uncertainties in its decay modes and the corresponding energy release. Note added: While we are preparing this manuscript, some other preprints \\cite{su} appeared, where the constraints on the gravitino dark matter scenario are studied. We found that the studies in \\cite{su} are quite exhaustively, where the scenarios of neutralino NLSP, stau NLSP as well as sneutrino NLSP are all considered. Compared with the studies in \\cite{su}, the characteristic of our study is that we performed a scan over the mSUGRA parameter space and presented the allowed regions in terms of original mSUGRA parameters. In addition, the BBN constraints on the EM energy release are more stringent in our study since we required such energy release settle the discrepancy between the BBN prediction and the WMAP data for $^{7}{\\rm Li}$ abundance." }, "0405/astro-ph0405366_arXiv.txt": { "abstract": "In discussing the short{}-time variability of extragalactic X{}-ray sources, we focus on explaining how wavelet transform in conjunction with non{}-linear filtering methods are used to remove Poisson statistics and wobble{}-related variability from the data. This enables the resolution of their intrinsic stochastic, quasi{}-periodic and deterministic variable components, as well the determination of their persistence. As detailed examples we review and extend the application of these techniques to ROSAT data for the Seyfert I galaxy NGC 5548 and for the QSO 3C273. Besides illustrating the discriminating power of these methods, these treatments confirm the intrinsic character of the transient quasi{}-periodic and deterministic events in NGC 5548 and 3C273 and enable us to compare both the elementary events and the short{}-time variability in the two types of sources. These methods can be applied in the investigation of intrinsic source variability in other spectral regions. ", "introduction": "{ A common belief in X{}-ray astronomy is that the spacecraft wobble precludes the possibility of obtaining correct information about the variability of observed sources. \\par} { The spacecraft wobble was introduced in order to prevent a given source from always being detected by the same pixel/pixels of the sensor. The linear wobble in ROSAT has been replaced by 2{}-dimensional ``dithering'' in the recent Chandra satellite, but the purpose of the technique remains unchanged. In ROSAT a consequence of the wobble is that the mesh wires in the entrance window of the instrument occult the observed source, thus introducing an instrumental variability. Such instrumental variability complicates observations, but it is far from destroying the information about the real variability of the source. On the contrary, in different physical measurements the emission from the source is chopped in order to increase the possibility of detection. In nature, certain species, like jumping spiders (\\textit{Salticidae}) acquire visual data by sweeping an essentially linear retina back and forth perpendicularly to its larger dimension \\citep{land69}. It has been demonstrated that the oscillating retina not only increases the resolution of the image, but also improves the over{}-all performance of vision, including its temporal resolution. This vision{}-enhancing principle is planned for use in the next generation of Mars rovers (New Scientist, 2001). \\par} { Modern signal processing techniques offer a wide range of possibilities for separating different components in the signal. In the case of X{}-ray observations the recorded signal is, as a rule, polluted by Poisson statistics, due to small numbers of recorded photons. However, during recent years a wavelet technique has been developed \\citep{liszka99} which removes much of the Poisson statistics from the data. In the present work the question whether the spacecraft wobble obscures the intrinsic variability of the source will be addressed using recent information processing techniques. Data from two intense X{}-ray sources: Crab and NGC 5548 were used in the present study. \\par} { Variability is an important characteristic of many astronomical systems, and rapid, apparently random variability is distinctive of many galactic and extra{}-galactic objects which are thought to harbor black holes, or other types of compact objects \\ {}-{}- X{}-ray binaries, QSO's, AGN and sources of gamma{}-ray bursts.It has proved very difficult to analyze in any reliable way the various components of the variability data from these objects in order to learn more about their underlying causes within the sources themselves. Here we review and extend the application we have made elsewhere \\citep{liszka99,liszka00a,liszka00b} of wavelet transform along with non{}-linear filtering techniques to the rapid X{}-ray variability of Seyfert 1 (S1) galaxies and QSO's. In particular, we show how we have employed these methods \\ to remove the Poisson{}-statistical and wobble{}-related variability from the ROSAT data for the S1 NGC 5548 and for the QSO 3C273. In doing so have confirmed the presence of intrinsic transient low{}-level deterministic and even quasi{}-periodic structures in data, which strongly indicate the occurrence of quasi{}-regular elementary events in the sources themselves. \\par} { These results were first reported in \\citet{liszka00a}. This paper presents a more detailed, careful and improved treatment of the methods used and the results obtained, particularly of the wobble{}-related and Poisson{}-statistical influences which must be removed to facilitate the determination of the properties of these sources. Our work confirms and strengthens observationally based support for the intrinsic character of the quasi{}-periodic and deterministic events we first reported in \\citet{liszka00b}, and significantly extends it to determination of the persistence of such events, their time{}-scales, and amplitudes. The methods we describe and apply here can obviously be applied to a wide{}-range of variable astronomical phenomena. Further confirmation of these results, of course, is crucial {}-{}- particularly by using data from more recent X{}-ray satellites, such as Chandra. \\par} \\begin{figure}[tb] \\center \\includegraphics[height=.57\\hsize]{fig1} \\figcaption{Formation of the observed photon train\\label{fig1}} \\end{figure} { The variability spectrum of an X{}-ray source may be a combination of deterministic and stochastic components (see the upper part of Fig.~\\ref{fig1}). The combined photon flux is increasingly affected by the Poisson statistics with the increased distance from the source. Finally, the photon train, measured at the spacecraft, is modulated by the wobble in a way which depends both on the apparent luminosity and on the spatial properties of the source. \\par} { Efficient deconvolution of the Poisson statistics from the time{}-scale spectra (frequency domain) became possible after of the introduction the wavelet transform into studies of photon trains \\citep{liszka99}. Using the wavelet transform as a pre{}-processing tool, non{}-linear modelling of the photon train \\ variability may be performed using a neural network technique \\citep{liszka00a,liszka03}. The problem of deconvolution of the spacecraft wobble may also be properly addressed using this technique. \\par} ", "conclusions": "{ We have reviewed and discussed some of the principal obstacles to reducing variability data from astronomical sources {}-{}- including Poisson statistics and instrumental motion, for example satellite wobble {}-{}- and the powerful methods of wavelet transform spectral analysis and non{}-linear filtering for overcoming these. We illustrated their application in a detailed way by applying these techniques to the rapidly variable X{}-ray emission from the Seyfert 1 galaxy NGC 5548 and to the QSO 3C273, confirming the presence and specifying some of the transient quasi{}-periodic and deterministic intrinsic events occurring in these sources. \\par} { In particular, after removing the modelled Poisson statistical and wobble{}-related variability components from the ROSAT time{}-scale spectra of 30 observation periods of NGC 5548 X{}-ray data, we found that there is significant residual variability. That is formed by a number of elementary events with a well defined time scale and magnitude. The elementary events seem to be distributed over the entire analyzed range of the time scale spectrum, both in time scale and magnitude domains. There is an indication of increasing intensity towards small magnitudes. Low wobble harmonics suppress, rather than enhance, the observed elementary events. It has been found by simulation that a wobble harmonic efficiently obscure other spectral peaks deviating less than 4\\% of its time scale. Typical elementary events last at least 3 basic time scales. The process responsible for the generation of elementary events has an internal bandwidth of 6\\% (allowed deviations of the time scale). \\par} { For a QSO type source (3C273) the elementary events are confined within a narrow range of normalized wavelet coefficient magnitudes around 10\\%. The elementary events in 3C273 are more intense, which may indicate that they are more persistent than the events observed in NGC 5548. The apparent compression of the events around 10\\% of the normalized wavelet coefficient magnitude could indicate the existence of a stronger, intrinsic stochastic component in the QSO type of source. However, a study of modelled variability spectra for a larger number of sources indicates that if there is an excess of an intrinsic stochastic component in the QSO sources, it must be less than a factor of 2 (see Fig.~\\ref{fig14}). \\par} { These methods are widely applicable to all sorts of variable data, most notably to astronomical source variability in other spectral ranges. Certainly, the results presented here need to be further confirmed by applying this analysis to data from more recent X{}-ray satellites, most notably Chandra, and to a larger number of other extragalactic sources. \\par}" }, "0405/astro-ph0405150_arXiv.txt": { "abstract": "We present results of the evolution of AGB stars of 3\\ms\\ and 5\\ms\\ with solar metallicity calculated with the Eggleton stellar evolution code ({\\sc{stars}}), which has a fully implicit and simultaneous method for solving for the stellar structure, convective mixing and nuclear burning. We introduce the concept of a viscous mesh in order to improve the numerical stability of the calculations. For the 5\\ms\\ star, we evolve through 25 thermal pulses and their associated third dredge-up events. We obtain a maximum helium luminosity of $1.7\\times10^9$\\ls\\ and significantly deep dredge-up after the second pulse. Strong hot-bottom burning is observed after the $5^{\\mathrm{th}}$ pulse. The 3\\ms\\ model is evolved through 20 thermal pulse events and we find third dredge-up after the 7th pulse. During the 14th pulse sufficient carbon has been brought to the surface to produce a carbon star. We find that dredge-up and the transformation into a carbon star occur at significantly smaller core masses (0.584\\ms\\ and 0.608\\ms, respectively) than in previous calculations for 3\\ms. ", "introduction": "The observed enhancements of carbon and $s$-process elements in asymptotic giant branch (AGB) stars show that material, processed by helium burning during thermal pulses, is being dredged up to the surface. Although qualitatively well understood, the extent and efficiency of this third dredge-up (TDUP) process has been very hard to model quantitatively, and remains one of the major uncertainties in AGB evolution. The amount of TDUP obtained for a star of a certain mass and composition differs greatly between one model calculation and another. It depends on many factors, such as the numerical treatment of convective boundaries \\citep{Lat89}, whether allowance for extra mixing is made \\citep{Hol88}, the mixing-length parameter \\citep{BSIV88} and the numerical resolution of the mass grid and time step \\citep{Stran97}. Recently, very efficient dredge-up for stars with small core masses has been found by various authors (Herwig et al.\\ 1997, Mowlavi 1999, Herwig 2000) if a diffusive form of convective overshooting is taken into account. In the calculation of the evolution of stars of $1-8$\\ms\\ through the thermally pulsing asymptotic giant branch (TP-AGB) phase of their evolution most evolutionary codes use only a partially self-consistent approach. The treatment of structure, mixing and nucleosynthesis is not simultaneous, with mixing often being solved in a separate iteration step from the other two. Examples include the codes of \\citet{Stran97}, \\citet{WG98}, \\citet{Her00} and \\citet{Amand02}. As the phenomenon of third dredge-up depends critically on the treatment of convection within a stellar structure calculation it is desirable to combine the calculation of nucleosynthesis, mixing and structure into a single, simultaneous step. The first attempt to do this was made by \\citet{Pols01}. In their calculations a 5\\ms\\ star was evolved without mass loss through the first 6 thermal pulses and deep dredge-up occured. The occurence of deep dredge-up is important for the formation of carbon stars. These are defined as stars that are M-type and have surface carbon-to-oxygen abundance ratios exceeding unity. Observational evidence suggests that these stars are of low mass, most likely between 1 and 3\\ms\\ (Iben 1981). However, there has been considerable difficulty in producing detailed theoretical models of carbon stars with low enough masses and luminosities. Of the early work on the subject, \\citet{BSIV88} were able to produce two carbon star models from initial masses of 1.2 and 2.0\\ms\\ under metal-poor conditions ($Z=0.001$), while \\citet{Lat89} produced a model initially of $M=1.5$\\ms\\ with $Z=0.02$. More recently low-mass carbon star models have been produced by \\citet{Stran97} and, with the aid of convective overshooting, \\citet{Her00}. However, the core masses found in these models are still too large to explain the carbon star luminosity function (Izzard \\& Tout 2004). While the approach of Pols \\& Tout enabled them to make significant progress with a fully implicit and simultaneous treatment of TP-AGB calculations, problems with numerical stability prevented them from evolving their model further than a few thermal pulses. In this work we have developed methods that overcome some of the numerical instability that occurs in a fully implicit treatment of the evolution. These methods are described in section 2. Results of the calculations for a 5\\ms\\ star are presented in section 3 and compared to other available models in section 4. In section 5 we present results of the calculations for a 3\\ms\\ and detailed comparisons are made to other models in section 6. ", "conclusions": "We have made a fully implicit calculation, with simultaneous solution of the structure, mixing and nuclear burning, of the evolution of both 5\\ms\\ and 3\\ms\\ stars. We have evolved the 5\\ms\\ star through 25 thermal pulses and their associated third dredge-up events. We find much stronger peak helium luminosities than are obtained by non-simultaneous codes and more efficient third dredge-up. Hot-bottom burning is found to occur during the interpulse. This prevents the build-up of carbon at the surface and leads to nitrogen being the most abundant CNO element. We have followed the evolution of a 3\\ms\\ star through 20 thermal pulses. We find deep third dredge-up to occur after the 7$^\\mathrm{th}$ thermal pulse and the star becomes a carbon star after the 14$^\\mathrm{th}$ thermal pulse. We find that these events occur at much lower core masses than those found by calculations with non-simultaneous codes. We also find much more efficient third dredge-up. The presence of efficient dredge-up and low core mass may help to explain the carbon star luminosity function. Working with the results of \\citet{Amand02}, Izzard \\& Tout (2004) demonstrated that, in order to fit the observed LMC and SMC carbon star luminousity functions, the minimum core mass for dredge-up to occur must be $0.07$\\ms\\ lower than those calculated with the MSSSP. While our models are at higher metallicity, we find dredge-up to occur when the core mass is $0.584$\\ms\\ or $0.05$\\ms\\ lower than the corresponding model of \\citet{Amand02}. We conclude that our calculations yield results that differ in very important ways from calculations made with non-simultaneous codes. We find third dredge-up to be more efficient and that it occurs at lower core masses. This may be due to the simultaneous solution of the equations although there are other minor differences between the numerical and physical details of different evolution codes. It should be noted that we have not conisdered the effects of mass loss in this work. Third dredge-up depends on the envelope mass (see e.g. Straniero et al. 2003) and the inclusion of mass loss would be expect to reduce the efficiency of third dredge-up." }, "0405/astro-ph0405199_arXiv.txt": { "abstract": "We report the discovery of several optical burst-like events from the low-mass X-ray binary MS~1603.6+2600 (UW~CrB). The events last for a few tens of seconds, exhibit a very fast rise and slow decay, and involve optical brightening of a factor of 2--3. The flares appear distinct from the lower level flickering and instead strongly resemble reprocessed type-I X-ray bursts as seen in a number of other neutron star low-mass X-ray binaries. In conjunction with the previously reported candidate X-ray burst, these confirm that the compact object in UW~CrB is a neutron star. We examine the optical burst brightness and recurrence times and discuss how the nature of the system can be constrained. We conclude that the source is most likely an accretion disk corona source at an intermediate distance, rather than a nearby quiescent system or very distant dipper. ", "introduction": "\\label{IntroSection} The X-ray source MS\\,1603.6+2600 was discovered in the Einstein Extended Medium Sensitivity Survey \\citep{Gioia:1990a} and associated with a faint ($R=19.4$) optical counterpart designated UW~CrB \\citep{Morris:1990a}. Its nature has remained a puzzle. \\citet{Morris:1990a} found the counterpart to be an eclipsing binary with an orbital period of 111.04\\,min and considered the source to be either a cataclysmic variable or low-mass X-ray binary (LMXB) hosting a neutron star. The emission line spectrum and optical to X-ray flux ratio favored the LMXB interpretation, with an accretion disk corona (ADC) source most likely. The implied distance was large, 30--80\\,kpc, making this high latitude source a halo object. \\citet{Hakala:1998a} reconsidered these possibilities and proposed another alternative -- a quiescent low-mass X-ray binary, likely a black hole system, which is much closer to us. An important clue was subsequently provided by \\citet{Mukai:2001a} who identified a strong X-ray flare in {\\it ASCA} data. While this appeared to resemble a type-I X-ray burst, the authors did not consider this identification conclusive. If the event was a type-I burst then it was faint, indicating either a very distant object in the halo, or an ADC source. Based on the X-ray lightcurve, \\citet{Mukai:2001a} favored the former of these interpretations, arguing that the source is a dipper rather than an ADC source. Finally, \\citet{Jonker:2003a} reported new {\\it Chandra} observations of the source, and favored the ADC interpretation, although they allowed that a quiescent system was still possible if the earlier X-ray flare was not a type-I burst. They rejected the distant dipper scenario, arguing that the optical luminosity would then be too high for a compact 2\\,hr binary. From a more theoretical standpoint, \\citet{Ergma:1993a} considered several evolutionary scenarios for the LMXB case, including degenerate and non-degenerate hydrogen rich mass donors and evolved helium stars. Again, bursts could be a crucial diagnostic. The presence, recurrence time, and duration of bursts can discriminate between systems with different mass transfer rates (e.g., the degenerate and non-degenerate cases discussed by \\citealt{Ergma:1993a}), and the burst properties will be sensitive to the chemical composition of the accreted material. To date, the only published X-ray burst from this source was that reported by \\citet{Mukai:2001a}, and this only yielded 60\\,counts. It is possible to also search for bursts in the optical, as the optical counterpart, while faint, is accessible to rapid photometry. Type-I X-ray bursts are expected to be manifested in the optical via reprocessed X-ray emission. This behavior has been widely seen in many other LMXBs for several decades (e.g., \\citealt{Grindlay:1978a} and many subsequent works). Optical bursts are dramatic, involving a brightening of a factor $\\sim 2$. Ultraviolet bursts are also present and are even more dramatic \\citep{Hynes:2004a}. We report here rapid optical photometry of UW~CrB. The primary goal of the program was to resolve the flickering contamination of the orbital variability, and this study will be presented separately. However, several optical bursts were serendipitously discovered, and we discuss those here. ", "conclusions": "We have reported the discovery of several resolved optical bursts from UW~CrB. These are almost certainly reprocessed type-I X-ray bursts, clarifying several characteristics of the source. i) For type-I bursts, the compact object must be a neutron star rather than a black hole or white dwarf. ii) The burst rate is relatively high, indicating an active rather than quiescent system, and thus a distance greater than a few kpc. iii) The optical burst flux is comparable to the similar source GS~1826--24, suggesting a comparable distance ($\\la 10$\\,kpc). It is thus most likely that UW~CrB is an ADC source (as also argued by \\citealt{Jonker:2003a}) rather than a distant dipper. Given its high Galactic latitude and intermediate distance, however, it must still be situated in the Galactic halo." }, "0405/gr-qc0405117_arXiv.txt": { "abstract": "We introduce a self-consistent stochastic coarse-graining method, which includes both metric and scalar field fluctuations, to investigate the back reaction of long wavelength perturbations in single-scalar driven inflation, up to the second (one loop) order. We demonstrate that, although back reaction cannot be significant during the last 70 e-foldings of inflation with a smooth potential, there exist non-smooth inflaton potentials which allow significant back reaction, and are also consistent with cosmological observations. Such non-smooth potentials may lead to the generation of massive primordial black holes, which could be further used to constrain/verify these models. ", "introduction": "In the context of inflationary cosmology, initial seeds of today's structures are generated in the de-Sitter phase of an inflationary universe, as a consequence of quantum vacuum fluctuations (see e.g. \\cite{MFB} for a comprehensive overview of the theory of cosmological fluctuations and \\cite{B03} for a recent introductory overview). The standard way of analyzing these fluctuations is through linear perturbation theory. However, the linear analyses is limited by the non-linear nature of the Einstein equations, and in particular, the presence of perturbations is likely to affect the evolution of the background cosmology at the non-linear level. The back reaction of short wavelength gravity waves on an expanding Friedmann-Robertson-Walker (FRW) cosmology is a well-understood problem \\cite{BHI}. However, in the context of inflationary and post-inflationary cosmology, the scalar metric fluctuations (fluctuations coupled to energy density and pressure perturbations) are believed to dominate over the effects of gravity waves, and thus, it is of great interest to understand the possible back reaction of these scalar perturbations on the cosmological background. Furthermore, during inflation, the phase space of infrared modes (defined as modes with wavelength greater than the Hubble radius) grows exponentially, whereas the phase space of the ultraviolet modes does not grow and since the amplitude of the associated metric fluctuations of these infrared modes does {\\it not} decrease in time , the back reaction of these infrared modes may grow to be significant. There have been different approaches to address this question. In \\cite{MAB,ABM}, the effective energy-momentum tensor formalism of \\cite{BHI} was generalized to study the back reaction of infrared modes on the spatially homogeneous component of the metric. The result was that, in a slow-roll inflationary background, the back reaction of infrared modes takes the form of a negative cosmological constant. This was later confirmed in \\cite{aw0} using very different techniques. The objections (see \\cite{unruh}) to these analyses are mainly concerned with the fact that \\cite{BHI,MAB,ABM} do not consider a local observable quantity, and even if the effect is not simply a gauge artifact, it is still not clear how a local observer could distinguish the back reaction of long-wavelength modes from an initial condition ambiguity in a homogenous universe. More importantly, calculating an ``observable'' from the spatially averaged metric will not in general give the same result as calculating the spatially averaged value of the observable. A later work by Abramo and Woodard in \\cite{aw2} tried to resolve this problem by identifying a local physical variable which describes expansion rate of the universe, then calculating the back reaction of cosmological perturbations on this quantity. A different approach was introduced in \\cite{GB1}, where a simple variable describing the local expansion rate was considered and calculated up to second order in the metric fluctuations in a model with a single inflation field \\footnote{This method has been recently used in \\cite{rasanen} to investigate the back reaction of UV modes in late Universe.}. In this approach, the leading contribution of back reaction of infrared modes to this variable was evaluated. When evaluated at a fixed value of the matter field $\\varphi$, the only physical clock available in this simple system, the dependence of the expansion rate on the clock time takes on exactly the same form as in an unperturbed background \\footnote{Back reaction of infrared modes for two scalar field models where one of the fields is the inflaton and the other one is the clock was investigated in \\cite{back2}.}. Thus, the leading infrared back reaction terms had no locally measurable effect in this system. Similar conclusions were reached in \\cite{aw3}. However, the fact that during inflation the phase space of infrared modes is growing in time, i.e. modes are exiting the Hubble radius, leads us to the conjecture that there might be a locally measurable effect just due to these crossings, and if we include the corrections due this effect, the results for single scalar field models may change. Since these modes are generated in the quantum mechanical ultraviolet limit, handling them in general is very difficult. Attempts, so far, in order to model the evolution through writing a classical Langevin-like equation for the coarse-grained scaler field, the so-called {\\it stochastic approach} \\cite{starobinsky}, fail to include the inflaton and metric fluctuations self-consistently \\footnote{A similar method has been used in \\cite{yokoyama} in a different context to investigate the behavior of weakly self-interacting scalar field in a de Sitter back ground.}. In this work, for the first time, we devise a self-consistent stochastic framework to study the evolution of the coarse-grained inflaton. While the approach uses the results of linear (metric+inflaton) perturbation theory for the ultraviolet and transitional modes, it is rigorously extended to the {\\it non-linear} infrared regime, which allows a consistent study of the evolution of fluctuations through their entire history. We then apply this framework to the problem of back reaction in inflationary cosmology, and investigate how/when the non-linear effects may change properties/predictions of the standard slow-roll inflation. The structure of the paper is as follows: In Section II, we explain how the linear theory of perturbations can be self-consistently generalized to the non-linear infrared regime, including both inflaton and metric fluctuations. We also introduce $\\Theta$, the local expansion rate, and how it is related to metric perturbations. In Sec. III, we show that the quantum generation of fluctuations can be modelled as a stochastic term in the equation of motion for $\\Theta$, coarse-grained over local Hubble patches. In Sec. IV, we apply this equation to a chaotic inflationary scenario with a quadratic potential $V(\\varphi)={1\\over2}m^2\\varphi^2$ and evaluate the local expansion rate with back reaction corrections due to infrared modes. Sec. V contains the generalization of the result from Sec. IV to any power-law potential. In Sec. VI, we investigate if back reaction can ever significantly change the inflationary predictions for our observable universe. Finally, Sec. VII summarizes our results and concludes the paper. Throughout this paper, we use the natural units, and set the Planck mass equal to unity, i.e. $8\\pi G = c= \\hbar =1$. ", "conclusions": "We use a stochastic coarse-graining method to approach the problem of infrared back reaction in a single scalar field inflationary cosmology. Using the coarse-grained generalized Bardeen parameter, we self-consistently derive the stochastic Langevin equation that governs the evolution of the local expansion rate. Solving this equation to second order (one loop) for chaotic inflation with a power law potential, we see that there is a non-vanishing average back reaction term due to the infrared modes that have left the Hubble radius. Although, the size of this effect will be negligible (within the observable universe) for smooth inflaton potentials, we show that it is possible to have non-smooth potentials which allow significant back reaction happening on sub-horizon scales, without breaking any observational constraints. A possible implication of such non-smooth potentials is the generation of a population of massive primordial black holes, which could be further used to constrain or verify these models. \\medskip \\centerline" }, "0405/astro-ph0405220_arXiv.txt": { "abstract": "We present an algorithm for parallelising the TreePM code. We use both functional and domain decompositions. Functional decomposition is used to separate the computation of long range and short range forces, as well as the task of coordinating communications between different components. Short range force calculation is time consuming and benefits from the use of domain decomposition. We have tested the code on a Linux cluster. We get a speedup of $31.4$ for $128^3$ particle simulation on $33$ processors; speedup being better for larger simulations. The time taken for one time step per particle is $6.5\\mu$s for a $256^3$ particle simulation on $65$ processors, thus a simulation that runs for $4000$ time steps takes $5$ days on this cluster. ", "introduction": "Observations of large scale structures like galaxies, clusters of galaxies along with observations of the cosmic microwave background radiation (CMBR) can be put together in a consistent framework if we assume that the large scale structures formed by gravitational amplification of density perturbations \\citep{tp93,peebles,peacock,lss_review}. These perturbations had a very small amplitude at the time of decoupling of matter and radiation, hence the highly isotropic character of the CMBR. Perturbations grow as overdense regions accrete mass and galaxies form when such regions are dense enough for star formation to take place. Early evolution of perturbations can be studied analytically using perturbation theory and approximation schemes. A detailed study of non-linear evolution of density perturbations requires the use of numerical simulations. Several methods have been developed for simulating gravitational clustering and formation of large scale structures, e.g. see \\citet{Bertsc98} for a review. The main driving force for these developments has been the need to simulate large systems in great detail while keeping errors in control. The emergence of Beowulf clusters as an affordable platform for high performance computing has given a fresh impetus to this activity, and the focus has shifted to algorithms that can be parallelised easily on such platforms \\citep{thesis,tpm,treepar96,tpmn,Vspr,mlapm,tpmnew,partreepm,gotpm,pmfast}. In this paper we present an algorithm for a parallel TreePM code. The TreePM method~\\citep{treepm,error_treepm} combines the tree code~\\citep{bh86} with a Particle-Mesh (PM) code, e.g. see \\citet{jbtp_pm,sim_book}. A brief summary of the TreePM method is given below, we refer the reader to \\citet{treepm}, and, \\citet{error_treepm} for more details and comparison with similar methods. Description of the TreePM method is followed by a discussion of the parallelisms inherent in the algorithm. In later sections we proceed to discuss our implementation and the performance. ", "conclusions": "We have presented an algorithm for parallelising the TreePM code on a Beowulf cluster. This code has been verified by comparing the final positions and velocities of particles in some test cases with the output of the sequential code, therefore the error profile of this code is same as the sequential TreePM code \\cite{error_treepm}. Even though we have tended to optimise the CPU time required at the cost of memory requirements, the maximum memory requirement per node is about $80$ bytes times the number of particles for the double precision code. We need up to $160$~MB per node for $128^3$ simulations and $1.25$~GB per node for $256^3$ simulations. These numbers represent the maximum memory requirements and for much of the time memory requirement is much smaller than this. Memory requirements can be reduced by about $25\\%$ by reorganising the code and adding a master node to gather positions and velocities of particles from nodes that are calculating the short range force. For $128^3$ simulations we get a speedup of $31.4$ on $33$ processors and $39$ on $65$ processors. The time taken for one time step per particle is $6.5\\mu$s for a $256^3$ particle simulation on $65$ processors, thus a simulation that runs for $4000$ time steps takes $5$ days on this cluster. These results are for a simulation with a global time step and further optimisations in terms of individual time steps is being carried out. The GOTPM code \\cite{gotpm} has a better performance in terms of time taken per particle per step. Part of the speedup is due to use of a larger mesh for the long range force calculation, and the remainder is due to a much smaller $r_{cut}$ and a more relaxed cell acceptance criterion for calculation of the short range force. The results for speedup efficiency and wall clock time per particle compare well with the published numbers for other parallel N-Body simulation codes of this class, e.g., \\citet{Vspr,tpmnew}." }, "0405/astro-ph0405016_arXiv.txt": { "abstract": "Enhancement and spread of helium among globular cluster stars have been recently suggested as a way to explain the horizontal branch blue tails, in those clusters which show a primordial spread in the abundances of CNO and other elements involved in advanced CNO burning. Helium enhancement is unavoidable, if the matter responsible for the abundance spreads is identified with the matter lost by massive asymptotic giant branch stars, which evolve during the early phases of globular cluster life (D'Antona et al.\\ 2002). In this paper we examine the implications of the hypothesis that, in many globular clusters, stars were born in two separate events: an initial burst (first generation), which gives origin to probably all high and intermediate mass stars and to a fraction of the cluster stars observed today, and a second, prolonged star formation phase (second generation) in which stars form directly from the ejecta of the intermediate mass stars of the first generation. In particular, we consider in detail the morphology of the horizontal branch in NGC 2808 and argue that it unveils the early cluster evolution, from the birth of the first star generation to the end of the second phase of star formation. This framework provides a feasible interpretation for the still unexplained dichotomy of NGC 2808 horizontal branch, attributing the lack of stars in the RR Lyr region to the gap in the helium content between the red clump, whose stars are considered to belong to the first stellar generation and have primordial helium, and the blue side of the horizontal branch, whose {\\it minimum} helium content reflects the helium abundance in the smallest mass ($\\sim 4$\\Msun) contributing to the second stellar generation. This scenario provides constraints on the required Initial Mass Function, in a way that a great deal of remnant neutron stars and stellar mass black holes might have been produced. ", "introduction": "Chemical inhomogeneities in globular cluster (GC) stars are giving crucial clues on the formation and evolution of this most ancient stellar population. Self-- enrichment mechanisms are being examined as ones of the possible -- and most plausible -- causes of abundance spread in many GCs. \\\\ The question of self--enrichment in GCs was raised for the first time in connection with the wide spread in metallicity observed in $\\omega$Cen \\citep{freeman-rodgers1975, freeman-norris1981, mallia-pagel1981}. Recently there has been a large amount of work on this problem (f.e., Smith et al.\\ 2000), an aspect of which is the ability of a cluster to retain the matter lost by asymptotic giant branch (AGB) stars. \\citet{gnedin2002} observe that $\\omega$Cen is not special among GCs in this respect, and that a few dozen other clusters should be capable of doing so. In fact, contrary to a substantial uniformity of abundance of heavy elements (apart from the quoted $\\omega$Cen), light elements that are susceptible to abundance changes from proton-capture reactions, such as the pp, CN, ON, NeNa, and MgAl cycles, exhibit star-to-star abundance variations in many GCs, far in excess of the modest variations seen in halo field stars --see, e.g., \\citet{smi87}, Kraft (1994)\\nocite{kra94}, and Sneden (1999,2000)\\nocite{sne99,sne00}. In recent years, observations of these abundance spreads at the turnoff and among the subgiant stars \\citep[e.g.,][]{gratton2001} have shown that these anomalies must be attributed indeed to some process of self--enrichment occurring at the first stages of the life of the cluster, starting as soon as all the supernovae have already exploded (expelling from the clusters their high velocity ejecta) and the massive AGB stars begin to evolve. At an epoch starting $\\sim 5 10^7$yr from the birth of the first stellar generation, the massive AGBs lose mass through low velocity winds, so that it can be reasonably speculated that these winds remain inside the cluster. \\\\ The massive AGB envelopes are the ideal place to manifacture elements through nuclear reactions in which proton captures are involved, as they are subject to hot bottom burning (HBB) \\citep[e.g.][]{ventura2001, ventura2002}, although a quantitative reproduction of the observed abundance spreads, e.g., of the oxygen vs. sodium anticorrelation \\citep{denis2003, ventura2004}, is still far from being available. One important feature of this hypothesis is that the chemistry of the stars which are formed in the second generation is not {\\it random}, but reflects the continuous variation of HBB conditions with the decrease of the progenitor evolving mass, until the star formation process ends \\citep[see the discussion in][]{dantona2003}. Thus the ejecta of these stars can be the source of the star-to-star abundance variations.\\\\ The hypothesis that the ejecta are accreted on the already formed stars \\citep{dgc1983, thoul2002} has been recently falsified by several observations --see, for a summary of these problems, \\citet{gratt-annualrev}. In this paper we further explore the hypothesis that such ejecta continuously form second generation stars for a time lasting about 200Myr \\citep{dantona2002}. \\\\ We argue that NGC 2808 is one of the clusters able to retain the AGB ejecta, even if to a degree lower than $\\omega$Cen, which is capable of retaining at least part of SNII ejecta (f.e., Smith et al.\\ 2000). Large variations in Na abundance have been observed among the stars of this cluster \\citep{carretta2003}, at all luminosity levels along the red giant (RG) branch. So members of this cluster have suffered pollution by Na--enriched matter, as observed in M13, M5, M4, M92, M15, NGC 6752. We do not have information on main sequence abundances, as in NGC 6752 \\citep{gratton2001} where a primordial origin for chemical inhomogeneities appears the only possible explanation, nevertheless the hypothesis of self-- enrichment in NGC 2808 is supported by the observed Na-- spread. We expect the most massive and concentrated clusters to be in the best conditions to retain mass lost by cluster members. NGC 2808 has a central density of log ($\\rho$) = 4.9 (\\Msun/pc$^3$, Pryor \\& Meylan 1993), among the largest densities found in Galactic GCs; the total visual magnitude is \\Mv\\ = -- 9.26 \\citep{harris1996}, which, assuming a ratio $M$/L$_{\\rm v}$ = 3, gives a total cluster mass well above $10^6$ \\Msun. According to Gnedin et al.\\ (2002, Fig.\\ 2), NGC 2808 would be at present quite able of retaining stellar winds of about 15 km s$^{-1}$, the characteristic terminal velocity of AGB winds (Loup et al.\\ 1993). Since according to \\citet{gnedin-ostriker1997} the destruction time for this cluster is about 6 Hubble times, present conditions should be close to those at formation epoch.\\\\ The hypothesis that GC stars were formed in two different star formation events was intended to explain the presence of long blue tails in the HBs \\citep{dantona2002} in GCs with large abundance spread \\citep{cat1995}. We have shown that blue tails can be the result of the evolution of main sequence stars with larger than average (cosmological) helium abundance. These latter stars have a mass smaller than average at the giant tip at present age and, for similar mass loss along the RG branch, will have a smaller mass (and bluer location) on the horizontal branch (HB). Helium enrichment is a natural outcome of the AGB evolution \\citep{ventura2002}, together with other chemical anomalies which can help to explain the observed abundance spreads. We leave aside the question of whether the winds from intermediate mass stars (IMS) can collect in the cluster central region, and here give origin to a new stellar generation, and concentrate on two other important issues: \\begin{enumerate} \\item we have to specify which initial mass function (IMF) is needed to give origin to enough stars in this second formation phase; \\item we have to explore more in detail which HB morphologies may derive from this process. In particular, the second star formation process may end when the helium content in the last ejecta (which give birth to stars) is still larger than the primordial helium content. This will produce {\\it a gap in the helium content}, between the original population of low mass stars, and the last born low mass stars. Does such a gap produce visible consequences in the distribution of stars in the HR diagram? \\end{enumerate} In this paper we attack these two problems, by extending the approach by \\citet{dantona2002} to a cluster having not only a blue HB with a long blue tail, but also a red HB. We will examine the case of the cluster NGC 2808, whose HB stellar distribution shows a well populated red clump, an extended blue HB, but it contains only two RR Lyr stars. This cluster has been studied intensively in recent years. The only way to reproduce the HB distribution is apparently to invoke a bimodal (or possibly multimodal) mass distribution \\citep[e.g.][]{fusipecci-bellazzini1997, catelan1998}. Differences in age or metal content among cluster members should be of a size not supported by observations. Bedin et al.\\ (2000) make an accurate study of the reddening, differential reddening and metallicity spread in the cluster, and conclude that the $\\sigma$ on [Fe/H] is 0.03 on the Carretta \\& Gratton scale (1997) and 0.05 on the Zinn \\& West scale (1984). On the other hand, their CM diagram and the one by Walker (1999) leave little doubt on the matter. As for an age difference, it should be of about 2--4 Gyr as required by the jump in colour (and so in mass, about 0.04 \\Msun\\ at the cluster metallicity) across the empty RR Lyrae region: an occurrence which appears rather unlikely. Besides, as noted by Rosenberg, Recio- Blanco \\& Garcia-Marin (2004), NGC 2808 is coeval with clusters of similar metallicity but much shorter HBs such as NGC 362, NGC 1261, NGC 1851. Similarly, not very much of definitive can be said on the role of rotation on HB distribution. After the many recent investigations on the matter, Sweigart (2002) concludes that: ``the impact of mixing and rotation on the HB remains to be clearly established\". Therefore we think important to explore new models for this complex cluster.\\\\ We here explore the possibility that the dichotomy in the HB is due to the differing helium contents of the clump stars and blue HB stars. This assumption is certainly parametric and has no pretence of uniqueness, but it has three attractive features: i) we know from the RG branch that the cluster stars show a dispersion in the abundances linked to the hot CNO cycle; ii) if we accept that the massive AGB stars are responsible for this abundance dispersion, all AGB models predict that the AGB winds are helium enriched; iii) the epoch of star formation from the AGB winds must necessarily end when massive AGBs are still evolving, and this will naturally produce a helium content gap.\\\\ It is important to stress how the large amount of information on the chemical composition of the members of GCs has changed profoundly our perception of these stellar systems. The impossibility of a simple solution for their formation and evolution - one star generation, a common age, one common composition - obliges to look for more complex solutions, which have to consider the detailed history of star formation, element production and dynamical evolution, and their interactions. While we can make reasonable guesses on element production, we lack any real knowledge of the dynamical history of GCs. So in this investigation we relied on a possible chemical history of NGC 2808, with only conjectures on dynamical processes. ", "conclusions": "In this paper we have used the galactic GC NGC 2808 as a laboratory to experiment on recent ideas on the early evolution of GCs. The HB in this cluster shows a very peculiar morphology: it is well populated in the red clump and in the blue and extremely blue region, but almost devoid of stars in the RR Lyr region. This morphology can be understood in terms of the coexistence of two stellar generations, separated by a difference between the helium content of the first one and the minimum helium content of the second one ($\\Delta Y \\simeq 0.03$). The small luminosity difference between the red and blue side of the horizontal part of the HB is also explained by this helium difference. \\\\ The peculiar HB morphology of NGC 2808 is therefore consistent with the hypothesis that the stars in many GCs were formed in two different events, a hypothesis primarily invoked to explain the abundance spreads in the light elements involved in the hot CNO cycle \\citep{cottrell-dacosta, dantona2002}. Using an average relation between the progenitor AGB mass and the helium content of its ejecta and the HB data by \\citet{bedin2000}, we derive two points of the IMF of the cluster stars: 1) the point at M=0.82\\msun, that is the mass function value for the stars which today are found in the red clump, having the primordial helium content; 2) the point for the mass range $\\sim 4.1 - 5$\\msun , whose ejecta should populate the most populous clump of the blue HB. Although this second point is somewhat model dependent, it clearly results that we need many more AGB stars than predicted by an IMF similar to those inferred for many stellar environments \\citep{kroupa2001}. We favour then the possibility that many of the low mass stars of the first stellar generation have been lost, but a dynamical study of these first stages of star formation is needed. \\\\ We comment on the large number of NSs --and possibly also BHs-- which our model predict, and which would make easier to understand the population of millisecond pulsars in many GCs and the possible presence of intermediate mass BHs in the clusters M15 and NGC 6752. \\\\ As for possible checks of the proposed scenario, the disappearance of surface helium (through sedimentation) at temperatures at which it would become observable, makes direct measurements impossible. A way to check the multiple star generations hypothesis could be the accurate measurement of main sequence masses in double stars, as well as the detailed analysis of the behaviour of various CM diagram features, such as the magnitude difference between the RG branch bump and the turnoff." }, "0405/astro-ph0405293_arXiv.txt": { "abstract": "{ We present VLA data at 330 MHz and 1.5 GHz of the radio emission observed in the cooling flow cluster A2626. By producing images at different resolutions we found that the radio source consists of different components: an unresolved core plus a jet-like feature, two elongated parallel features, and an extended diffuse emission (radio mini-halo). Low resolution images allow us to derive morphological and spectral information of the diffuse emission: the radio mini-halo is extended on a scale comparable to that of the cooling flow region and is characterized by amorphous morphology, lack of polarized flux and very steep spectrum which steepens with distance from the center. We then applied to this new mini-halo source a model for particle re--acceleration in cooling flows (Gitti, Brunetti \\& Setti 2002). In particular, we found that its main radio properties (brightness profile, integrated radio spectrum and radial spectral steepening) can be accounted for by the synchrotron radiation from relic relativistic electrons in the cluster, which are efficiently re-accelerated by MHD turbulence amplified by the compression of the cluster magnetic field in the cooling flow region.} ", "introduction": "The cluster A2626 (z=0.0604) is a good candidate to study the interaction between the X--ray emitting intra--cluster medium (ICM) and radio emitting plasma in clusters of galaxies. This cluster hosts a relatively strong cooling flow (White, Jones \\& Forman 1997) and contains a very unusual radio source exhibiting a compact unresolved core and a diffuse structure (Roland et al. 1985; Burns 1990). Earlier radio observations showed that the compact component is associated with the centrally dominant elliptical galaxy (Owen, Ledlow, \\& Keel 1995), while the diffuse emission has no optical counterpart. Comparisons with X--ray data revealed an enhanced X--ray emission spatially coincident with the radio source, thus providing strong observational evidence for a connection between the hot, X--ray gas and the radio plasma. In addition, the X--ray map of the cooling flow region shows an elongation coincident with the diffuse radio component (Rizza et al. 2000). We studied in detail the radio properties of A2626 in order to better investigate the nature of the interaction between the ICM and the radio plasma and in particular the origin of the diffuse radio emission in the core of cooling flow clusters. A Hubble constant $\\mbox{H}_0 = 50 \\mbox{ km s}^{-1} \\mbox{ Mpc}^{-1}$ is assumed in this paper, therefore at the distance of A2626 $1'$ corresponds to $\\sim$ 95 kpc. The radio spectral index $\\alpha$ is defined such as $S_{\\nu} \\propto \\nu^{-\\alpha}$. ", "conclusions": "" }, "0405/astro-ph0405546_arXiv.txt": { "abstract": "{We present the catalogue of the REFLEX Cluster Survey providing information on the X-ray properties, redshifts, and some identification details of the clusters in the REFLEX sample. The catalogue describes a statistically complete X-ray flux-limited sample of 447 galaxy clusters above an X-ray flux of $3\\ 10^{-12}$ erg s$^{-1}$ cm$^{-2}$ (0.1 to 2.4 keV) in an area of 4.24 ster in the southern sky. The cluster candidates were first selected by their X-ray emission in the ROSAT-All Sky Survey and subsequently spectroscopically identified in the frame of an ESO key programme. Previously described tests have shown that the sample is more than 90\\% complete and there is a conservative upper limit of 9\\% on the fraction of clusters with a dominant X-ray contamination from AGN. In addition to the cluster catalogue we also describe the complete selection criteria as a function of the sky position and the conversion functions used to analyse the X-ray data. These are essential for the precise statistical analysis of the large-scale cluster distribution. This data set is at present the largest, statistically complete X-ray galaxy cluster sample. Together with these data set we also provide for the first time the full three-dimensional selection function. The sample forms the basis of several cosmological studies, one of the most important applications being the assessment of the statistics of the large-scale structure of the universe and the test of cosmological models. Part of these cosmological results have already been published. \\thanks{The full versions of Tables 2 through 9 will be available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/ as well as on our home page http://www.xray.mpe.mpg.de/theorie/REFLEX/DATA} } ", "introduction": "Clusters of galaxies are the largest building blocks of our Universe that can still reasonably well be characterized as unique objects. This makes them on one hand very important large-scale astrophysical laboratories in which a large variety of astrophysical processes can be studied in well characterized environments. For these laboratories we can measure for example their total gravitational mass, their matter composition, the internal gas density, temperature, and pressure of the intergalactic medium, their distance, and other important properties. The best basis for such astrophysical studies is a well documented catalogue of galaxy clusters to choose the best suited objects for the prospective study (e.g. B\\\"ohringer et al. 2001b). On the other hand X-ray selected galaxy clusters are very good tracers of the large-scale structure of the Universe. Since there is a quite well understood relation between the distribution of galaxy clusters with known mass and the dark matter distribution, the statistics of the large-scale matter distribution in the Universe can be derived from the distribution of clusters in a well selected, statistically complete sample. This study of the large-scale structure was the main objective for the construction of the REFLEX sample. Several results on the construction of the sample (B\\\"ohringer et al. 2001a, Paper I), the assessment of the large-scale structure (Collins et al. 2000, Paper II; Schuecker et al. 2001a, Paper III; B\\\"ohringer et al. 2002, Paper IV; Schuecker et al. 2002, 2003a (Papers VI and VII), 2003b, Kerscher et al. 2001), on the statistics of substructure in REFLEX clusters Schuecker et al. (2001b), on the statistics of the cluster galaxy velocity dispersions (Ortiz-Gil et al. 2003), and on the X-ray temperatures of the most luminous, distant REFLEX clusters (Zhang et al. 2003) have already been published. Several further papers are in preparation. Due to the close correlation of X-ray luminosity and mass for clusters of galaxies (e.g. Reiprich \\& B\\\"ohringer 2002) the detection and selection of the sample clusters is currently best performed through the cluster X-ray emission. The ROSAT All-Sky Survey (RASS), which is still the only all-sky or wide-angle X-ray survey performed with an imaging X-ray telescope, is by far the best basis for such cosmological studies. It has been used previously in several projects to construct statistical galaxy cluster samples (Pierre et al. 1994, Romer et al. 1994, Ebeling et al. 1996, 1998, 2000, Burns et al. 1996, De Grandi et al. 1999, Ledlow et al. 1999, B\\\"ohringer et al. 2000, Henry et al. 2001, Cruddace et al. 2002, 2003, Ebeling et al. 2001, 2002, Gioia et al. 2003). Part of these projects were studies connected to and profiting from the REFLEX survey program. None of the previous projects covers an area in the southern sky as large as REFLEX, except for the XBACS Abell cluster survey (Ebeling et al. 1996), which is shallower and restricted to those clusters previously identified by Abell (1958) and Abell, Corwin, and Olowin (1989). The REFLEX catalogue of 447 clusters provides presently the largest statistically complete X-ray cluster sample. The volume of the Universe that is probed is larger than that covered by any present galaxy redshift survey except for the Sloan Digital Sky Survey, which goes to a slightly larger depth but will only cover about half the sky area of that covered by REFLEX, when completed. The paper is organized as follows. In section 2 we describe the survey and the selection characteristics. Section 3 provides a brief description of the X-ray data reduction and section 4 describes the redshift determination and the cluster galaxy redshift statistics. The main catalogue is presented in section 5 and some of its properties are reviewed in section 6. In the latter section we also provide the numerical data and the recipe to construct the survey selection function in one and two dimensions for any flux limit equal to or above the nominal REFLEX flux limit. Section 7 gives some further information on the identification and the properties of some individual clusters. Section 8 lists close cluster pairs and clusters with double or multiple X-ray maxima found in the REFLEX catalogue, and we describe in more detail those clusters where multiple redshift clustering is observed in the line-of-sight of the X-ray source. In section 9 we compare the results with the previously derived survey samples and, finally, in section 10 we provide a summary and conclusions. Table \\ref{tab0} gives an overview of the information presented in this paper in tabular form. \\begin{table} \\caption{Overview on the data presented in this paper in tabular form} \\label{tab0} \\[ \\begin{array}{ll} \\hline \\noalign{\\smallskip} {\\rm Table} & {\\rm Content} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\ref{tab1} & {\\rm count~ rate~ to~ flux~ conversion~ (for~ z=0)} \\\\ \\ref{tab2} & {\\rm K-correction~ as~ a~ function~ of~ z~ and~ T_x} \\\\ \\ref{tab3} & {\\rm Flux~ conversion~ from~ the~ 0.1 - 2.4~ keV~ to~}\\\\ & {\\rm the~ 0.5~ to~ 2.0~ keV~ band} \\\\ \\ref{tab4} & {\\rm Flux~ conversion~ from~ the~ 0.1 - 2.4~ keV~ band~}\\\\ & {\\rm to~ bolometric~ flux} \\\\ \\ref{tab5} \\& \\ref{tab6} & {\\rm REFLEX~ cluster~ catalogue~ for~ h = 0.7~ and~ \\Lambda -cosmology} \\\\ \\ref{tab7} & {\\rm Further~ X-ray~ parameters~ of~ the~ REFLEX~ clusters} \\\\ \\ref{tab8} & {\\rm Sky~ coverage~ as~ a~ function~ of~ the~ flux~ limit} \\\\ \\ref{tab9} & {\\rm Angular~ modulation~ of~ the~ survey~ selection~ function} \\\\ \\ref{tab10} & {\\rm Close~ cluster~ pairs~ in~ the~ REFLEX~ sample} \\\\ \\ref{tab11} & {\\rm Clusters~ with~ multiple~ maxima~ in~ the~ REFLEX~ sample} \\\\ \\ref{tab12} & {\\rm Line-of-sight~ redshift~ clustering~ at~ the~}\\\\ & {\\rm position~ of~ REFLEX~ clusters} \\\\ \\noalign{\\smallskip} \\hline \\end{array}\\] \\end{table} The luminosities and other cluster parameters which depend on the distance scale are derived for a Hubble constant of $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$ and a cosmological model with $\\Omega_m = 0.3$ and $\\Omega_{\\Lambda} =0.7$ in the main tables of the paper. We also give in complementary tables provided only in electronic form the cluster properties for the previously most often used Einstein-de Sitter model with $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 1.0$ and $\\Omega_{\\Lambda} =0$ for an easier comparison with previous literature results. ", "conclusions": "" }, "0405/astro-ph0405636_arXiv.txt": { "abstract": "We use the Allen et al.~(2004) $Chandra$ measurements of x-ray gas mass fraction of 26 rich clusters to place constraints on the scalar-field dark energy model with inverse power law potential energy density. The constraints are consistent with, and typically more constraining than, those from other cosmological tests, and mildly favor the Einstein cosmological constant limit of the dark energy model. ", "introduction": "All indications are that the energy budget of our universe has recently come to be dominated by some form of dark energy, resulting in an accelerating cosmological expansion. This picture is supported by a number of measurements, including: Type~Ia supernova redshift-magnitude data (see, e.g., Wang \\& Mukherjee 2004; Nesseris \\& Perivolaropoulos 2004; Riess et al.~2004; Biesiada et al.~2004; Daly \\& Djorgovski 2004); cosmic microwave background (CMB) anisotropy data from the Wilkinson Microwave Anisotropy Probe (WMAP), with some input from other measurements (see, e.g., Bennett et al.~2003; Page et al.~2003); and, other measurements of CMB anisotropy, which indicate the universe is close to spatially flat (see, e.g., Podariu et al.~2001b; Durrer et al.~2003; Melchiorri \\& \\\"Odman 2003), in combination with continuing strong evidence for low non-relativistic matter density (Chen \\& Ratra 2003b and references therein). For reviews see Peebles \\& Ratra (2003), Padmanabhan (2003), Steinhardt (2003), Carroll (2004), and Sahni (2004). There are many proposed dark energy candidates.\\footnote{ For recent discussions of dark energy models and the observational situation see Bartolo et al.~(2003), Kratochvil et al.~(2003), Kaplinghat \\& Bridle (2003), Gong (2004), Liu (2004), Gorini et al.~(2004), Wetterich (2004), Matarrese et al.~(2004), Bludman (2004), Feng et al.~(2004) and references therein.} The original example of dark energy is Einstein's cosmological constant, $\\Lambda$, which has an energy density $\\rho_{\\Lambda}$ independent of time and space. The modern reincarnation of this is in the $\\Lambda$CDM model (Peebles 1984), where at low redshift non-relativistic matter, dominated by the also hypothetical cold dark matter (CDM), is the other major contributor to the cosmological energy budget. More recently, dark energy models in which the dark energy density varies only slowly with time and space have attracted much attention. A simple example of such a candidate is a scalar field ($\\phi$) with potential energy density $V(\\phi) \\propto \\phi^{-\\alpha}$, $\\alpha > 0$, at low redshift $z$ (Peebles \\& Ratra 1988; Ratra \\& Peebles 1988); in what follows we refer to this as the $\\phi$CDM model. The XCDM parametrization of varying dark energy density models approximates the dark energy by a fluid with a time-independent equation of state parameter $w = p/\\rho$, where $p$ is the pressure. The XCDM parameterization is accurate in the radiation and matter dominated epochs, but much less so in the scalar field dominated epoch when $w$ is time dependent (see, e.g., Ratra 1991). In the last two cases, consistent with observational indications, we consider only a spatially flat cosmological geometry; the $\\Lambda$CDM model considered here allows space curvature to be a free parameter. It is important to test dark energy models and constrain their parameters using as many techniques as possible. Different tests might provide different constraints on the parameters of the model, and comparison of results determined from different methods allow for consistency checks. A number of such cosmological tests have been developed. In addition to the Type~Ia supernova test mentioned above\\footnote{ The proposed SNAP/JDEM space mission to measure the redshift-magnitude relation to larger redshift should provide valuable data for constraining dark energy models (see, http://snap.lbl.gov/ and, e.g., Podariu et al.~2001a; Weller \\& Albrecht 2002; Rhodes et al.~2004; Virey et al.~2004).}, there has been discussion of the redshift-angular-size test (see, e.g., Zhu \\& Fujimoto 2002; Chen \\& Ratra 2003a; Podariu et al.~2003; Jain et al.~2003; Jackson 2003); the redshift-counts test (see, e.g., Loh \\& Spillar 1986; Newman \\& Davis 2000; Huterer \\& Turner 2001; Podariu \\& Ratra 2001); the strong gravitational lensing test (see, e.g., Fukugita et al.~1990; Turner 1990; Ratra \\& Quillen 1992; Waga \\& Frieman 2000; Chae 2003; Chae et al.~2004); and the redshift-expansion time test (see, e.g., Nolan et al.~2003; Alcaniz et al.~2003; Savage et al.~2004; Cepa 2004). Structure formation in time-variable dark energy models has also come under recent discussion (see, e.g., Mainini et al.~2003; Klypin et al.~2003; {\\L}okas et al.~2004; Mota \\& van de Bruck 2004), and CMB anisotropy data is providing useful constraints (see, e.g., Mukherjee et al.~2003a, 2003b; Caldwell \\& Doran 2004; Wang \\& Tegmark 2004; Jassal et al.~2004). In this paper, we use the x-ray gas mass fraction of rich clusters, as a function of redshift, to constrain the three simple dark energy models mentioned above. This test was introduced by Sasaki (1996) and Pen (1997), and further developed by Allen et al.~(2002, hereafter A02; also see Allen et al.~2003; Ettori et al.~2003; Allen et al.~2004, hereafter A04, and references therein).\\footnote{ This test builds on the zero-redshift cluster baryon mass fraction test discussed by White \\& Frenk (1991), Fabian (1991), White (1992), and White et al.~(1993).} The basic idea is as follows. Assuming that the rich clusters are large enough to provide a fair representation of the cosmological baryon and dark matter distributions (see, e.g., White 1992; White et al. 1993; Fukugita et al.~1998), the ratio of baryonic to total mass in clusters---the cluster baryon mass fraction---should be the same as the ratio of baryonic to non-relativistic mass in the whole universe---the cosmological baryon mass fraction. The zero-redshift cluster baryon fraction test allows for a determination of the non-relativistic matter mass density parameter $\\Omega_m$ from the measured cluster baryon fraction and an estimate of the baryonic mass density parameter $\\Omega_b$. If one focuses on the rich clusters (and not on those that might be in the process of collapsing), the cluster baryon mass fraction should be independent of the cluster redshift. The main contributors to the cluster baryon mass fraction are the cluster gas mass fraction and the cluster galaxy (stellar) mass fraction, with the cluster gas mass fraction dominating. The cluster gas mass fraction depends on the angular diameter distance (see, e.g., Sasaki 1996) so data on it as a function of redshift allows for another cosmological test: the correct cosmological model places the clusters at the right angular diameter distances to ensure that the gas mass fractions are independent of redshift. A02 and A04 focus on the x-ray emitting intracluster gas, and use the x-ray gas mass fraction of the total cluster mass, $f_{\\rm gas}$, to constrain cosmological parameters. Note that the optically luminous galaxy (stellar) mass in clusters is about $0.19 h^{0.5}$ times the x-ray emitting gas mass\\footnote{ Here $h$ is the Hubble constant in units of 100 km s$^{-1}$ Mpc$^{-1}$. The expression $0.19 h^{0.5}$ is from A02 and is based on Fukugita et al. (1998) who use a distance-independent stellar $M/L$, rather than a distance-dependent dynamical $M/L$, so it differs from the $0.19h^{1.5}$ used by White et al. (1993).}, so $\\Omega_b = \\Omega_m f_{\\rm gas} (1 + 0.19 h^{0.5})$. In $\\S$ 2 we summarize our method of computation. Results are presented and discussed in $\\S$ 3. We conclude in $\\S$ 4. ", "conclusions": "We use recent x-ray cluster gas mass fraction data from the $Chandra$ observatory to constrain cosmological parameters. These constraints are consistent with those derived from other cosmological tests for a range of parameter values in each of the three cases we consider, but mildly favor the spatially-flat $\\Lambda$CDM model. The $f_{\\rm gas}$ data is efficacious at constraining dark energy, and also provides a relatively tight and approximately model-independent constraint on $\\Omega_m$, $0.15 \\lesssim \\Omega_m \\lesssim 0.35$ at two standard deviations, which is in good accord with other recent estimates (Chen \\& Ratra 2003b; Bennett et al.~2003). Future $f_{\\rm gas}$ data should provide an even tighter constraint and is eagerly anticipated. \\bigskip We are indebted to S.~Allen for providing the cluster gas mass fraction data and helpful discussions. We also acknowledge helpful discussions with J.~Alcaniz, J.~Peebles, and Z.~Zhu, and support from NSF CAREER grant AST-9875031 and DOE EPSCoR grant DE-FG02-00ER45824. We thank the referee for useful advice." }, "0405/astro-ph0405492_arXiv.txt": { "abstract": "{ We investigated the MeV properties of 173 unidentified or only tentatively identified EGRET sources listed in the third EGRET catalogue, by analyzing the simultaneously collected COMPTEL MeV data for each individual source. The sources can generally be divided into 4 groups. In this paper we focus on one of these, a group of 22 EGRET sources for which we can provide additional constraining information: their spectral extrapolations from the energy range above 100 MeV towards lower energies overshoot the fluxes or upper limits derived simultaneously at MeV energies. This means that for these sources a spectral turnover/break between 1 MeV and 100 MeV is required. At least two of these sources, but most likely the majority of this sample, have the maxima of their gamma-ray luminosities in this energy band. The sources have rather soft EGRET spectra (average photon index $=$ 2.72$^{+0.08}_{-0.11}$), and seem to spatially cluster in the inner Galaxy. Variability analyses revealed 11 out of the 22 sources to be significantly variable. Object classes proposed as possible counterparts for the unidentified EGRET sources are discussed in the light of these additional constraints. ", "introduction": "One of the biggest mysteries left by the Compton Gamma-Ray Observatory (CGRO, 1991-2000) is that a large number of $\\gamma$-ray sources detected by the different CGRO experiments, in particular EGRET, still remain unidentfied. The EGRET experiment measured \\grays\\ above 30~MeV, most sensitively above 100~MeV. Out of the 271 sources listed in the third EGRET catalogue \\citep{Hartman99}, 170 are unidentified and 27 are only tentatively identified. Several classes of objects have been proposed as possible counterparts for those unidentified EGRET sources. Sources located at high galactic latitudes and being time variable are believed to be active galactic nuclei (AGN), in particular blazars. Sources, located at lower galactic latitudes, being steady and having low \\gray\\ fluxes, are found to coincide spatially with the Gould Belt \\citep{Gehrels00}. Some other low-latitude sources show positional correlations with supernova remnants (SNRs) and OB associations \\citep[e.g.][]{Romero99}. Steady sources with hard \\gray\\ spectra seem to be good candidates for young $\\gamma$-ray pulsars with ages of less than 10$^{6}$ years \\citep{Zhang00a}. Several sources (mainly located $|$b$|$ $<$ 10\\deg) might indicate a new class of $\\gamma$-ray emitting objects \\citep{Torres01}, because they do not coincide with any potential counterpart objects. The COMPTEL experiment aboard CGRO is sensitive to \\gray\\ photons between 0.75 and 30~MeV, thereby covering the softer \\gray\\ band adjacent to the EGRET one. Apart from transient $\\gamma$-ray bursts, unidentified \\gray\\ sources and AGN are the majority of the COMPTEL source detections. The first COMPTEL catalogue \\citep{Schonfelder00} lists 10 AGN and 9 unidentified $\\gamma$-ray sources; the sum of the rest (radio pulsars, stellar black-hole candidates, SNRs, and $\\gamma$-ray line sources) is about 12. Since COMPTEL and EGRET were mounted parallel on CGRO and both had a large field of view (the COMPTEL one being larger than the EGRET one), COMPTEL and EGRET observed simultaneously the sa\\-me sky region. To gain further knowledge on the unidentified EGRET sources, and to probe their nature, we analyzed the contemporaneous COMPTEL data on the unidentified EGRET sources to supplement the EGRET results. In this paper, we report the discovery of a subgroup of the unidentified EGRET sources whose \\gray\\ spectra are constrained by the MeV data: their spectral energy distributions have at least an MeV break but most likely an MeV peak. The paper is organized as follows: in Section 2 we briefly describe the COMPTEL instrument, the applied data analysis methods and the observational concept of CGRO, in Section 3 we present the analysis results and discuss them in Section 4. In Section 5, we finally present the conclusions. ", "conclusions": "By analyzing the contemporary COMPTEL observations of all unidentified or only tentatively identified EGRET sources, we found a subgroup of 22 sources for which we can provide spectral constraints for source modelling. Their spectra have to turn over between $\\sim$1 MeV and 100 MeV, and at least two of them, but most likely the majority, have their maximum luminosities somewhere in this energy band. Most of the sources are not detected by COMPTEL, however the simultaneously derived upper limits require the spectral bending. These sources have rather steep energy spectra in the EGRET band, and seem to be preferentially located in the inner galaxy, especially at low latitudes ($|$b$|$ $<$ 30$^{\\circ}$). Variability studies reveal that half of them are significantly variable above 100~MeV. Potential counterparts have to conform to these observational results. A blazar origin for the two high-latitude sources in this sample seems to be likely. Viable candidate counterparts for the steady low-latitude sources are: 1) young (age $<$ 10$^6$ years) pulsars, although the youngest with the strongest magnetic fields might be too strong in the X-ray domain, 2) old, recycled millisecond pulsars with a weak magnetic field (like PSR J0218+4232), and 3) SNRs and pulsar wind nebula whose synchrotron spectra are peaking at MeV energies (like the Crab nebula). For the variable low-latitude sources, XRBs, in particular microquasars/blazars by assuming a spectral analogy to the extragalactic objects, and isolated BHs would match the requirements. Case by case studies might reveal further insights in the nature of individual sources." }, "0405/astro-ph0405171_arXiv.txt": { "abstract": "Powerful radio galaxies at high redshift are highly useful in studies of early evolution of AGN-hosting galaxies because their observed optical and near infrared light are dominated by their stellar population rather than the nonthermal continuum emitted by the central engine of AGNs. In addition, the presence of AGN activity in them implies that a supermassive black hole has been already made in their nuclei. These properties allow us to investigate a possible starburst-AGN connection in early universe and then provide some crucial hints for the formation mechanism of supermassive black holes. Taking observational properties of high-$z$ powerful radio galaxies into account, we discuss a possible formation mechanism of supermassive black holes in their nuclei. ", "introduction": "The formation and evolution of supermassive black holes were once thought to be related only to some exotic phenomena such as nonthermal nuclear activity. Active galactic nuclei (AGNs), in particular, quasars and powerful radio galaxies, are basically minor populations among galactic nuclei. Therefore, it has long been presumed that the history of supermassive black holes (SMBHs) is independent from that of galaxies or that of stars. However, since the discovery of a massive dark object in the heart of Andromeda galaxy (M31; Dressler \\& Richstone 1988; Kormendy 1988), it has been reported that a number of nearby normal galaxies appear to have such a massive dark object in their nuclei (e.g., Kormendy \\& Richstone 1995). Then, in the 90s, beautiful evidence for a SMBH was found in one of famous nearby AGNs, NGC 4258, based on kinematic motion of H$_2$O masing clouds around the nucleus (Miyoshi et al. 1995). Further challenge made in this decade has brought the discovery of a tight correlation between the spheroidal mass and the SMBH mass (Magorrian et al. 1998; see also Silk \\& Rees 1998). Surprisingly, normal galaxies and AGN-hosting galaxies follow the same correlation (Gebhardt et al. 2000; Ferrarese et al. 2001; Wandel 2002). This suggests that most {\\it nucleated} galaxies have a massive dark object, a SMBH, despite of the nonthermal nuclear activity. It is also suggested that AGN phenomena are not associated with the formation process of a SMBH itself (e.g., Taniguchi 2002). More importantly, the formation of SMBHs may be physically linked to that of spheroidal component of galaxies. In order to explore this issue, it seems important to investigate young galaxies in early universe. For this purpose, we use observational properties of high-$z$ powerful radio galaxies\\footnote{The most distant HzPRG known to date is TN J0924$-$2201 (van Breugel et al. 1999). However, most HzPRGs are located at $z \\sim 2$ - 4 (e.g., McCarthy 1993).} (HzPRGs) because their host galaxies have been intensively investigated for these two decades (e.g., McCarthy 1993) and consider a possible scenario for the the formation of SMBHs at high redshift (see for review, Rees 1984; Haiman \\& Quataert 2004). ", "conclusions": "We summarize a possible scenario for the formation of a SMBH in nuclei of HzPRGs (as well as quasars) at $z \\sim 3$ (see also Table 1). \\begin{description} \\item{Step I.} A very massive star (i.e., a Pop III object) could born in a dark matter mini halo at $z \\sim 25$. Shortly after the formation, an IMBH with $M_\\bullet \\sim 100 M_\\odot$ could be left as a compact remnant; i.e., the first-stage seed black hole. \\item{Step II.} The gas accretion at the Eddington limit could grow up an IMBH to $M_\\bullet \\sim 10^7 M_\\odot$ at $z \\sim 8$ during the course of successive mergers among dark matter halos; this can be regarded as a seed SMBH. This growth takes 0.6 Gyr from $z \\sim 25$. \\item{Step III.} A few or several massive gaseous system with $M \\sim 10^{11} M_\\odot$ could merge into one. In their final phase, a large number of super star clusters could be born near the merger center because of strong dynamical action driven by binary of multiple seed SMBHs. \\item{Step IV.} Such super star clusters could sink into the merger center because of dynamical friction, leading to the formation of a SMBH with $M_\\bullet \\sim 10^9 M_\\odot $at $z \\sim 3$. This process takes $\\sim$ 1 Gyr, being shorter than the time interval between $z=8$ and $z=3$ given the currently accepted WMAP cosmology. \\end{description} The above scenario should be related to the formation of a massive giant elliptical galaxy. The very luminous starburst (or the initial starburst) could contribute to the formation of the spheroidal component. Therefore, the Magorrian relation could be established in the Step IV in this scenario. It is noted that such merger-driven ultraluminous starbursts could explain the Magorrian value of $M_\\bullet/M_{\\rm spheroid} \\sim 0.001$ (Bekki \\& Couch 2001). { \\scriptsize \\tablenum{1} \\tablewidth{6.5in} \\begin{deluxetable}{ccclc} \\tablecaption{% A Summary of the Proposed Scenario for the formation of SMBHs } \\tablehead{ \\colhead{$z$} & \\colhead{$T_{\\rm Univ}$ (Gyr)} & \\colhead{$\\Delta T$ (Gyr)} & \\colhead{Event} & \\colhead{$M_\\bullet$ ($M_\\odot$)} } \\startdata 25 & 0.35 & & Pop. III $\\Rightarrow$ IMBH & $\\sim 10^2$ \\nl & & 0.55 & gas accretion & \\nl 8 & 0.90 & & onset of mergers among massive halos & $\\sim 10^7$ \\nl & & 1.40 & starburst $\\Rightarrow$ super star clusters & \\nl 3 & 2.30 & & sinking into the merger center & $\\sim 10^9$ \\nl \\enddata \\end{deluxetable} }" }, "0405/astro-ph0405037_arXiv.txt": { "abstract": "Using a combination of data from the Antarctic Submillimeter Telescope and Remote Observatory (AST/RO), the Arizona Radio Observatory Kitt Peak 12m telescope and the Arizona Radio Observatory 10m Heinrich Hertz Telescope, we have studied the most active part of the R~CrA molecular cloud in multiple transitions of Carbon Monoxide, HCO$^+$ and 870\\micron\\ continuum emission. Since R~CrA is nearby (130 pc), we are able to obtain physical spatial resolution as high as 0.01pc over an area of 0.16 pc$^2$, with velocity resolution finer than 1 km/s. Mass estimates of the protostar driving the mm-wave emission derived from HCO$^+$, dust continuum emission and kinematic techniques point to a young, deeply embedded protostar of $\\sim$0.5-0.75 M$_\\odot$, with a gaseous envelope of similar mass. A molecular outflow is driven by this source that also contains at least 0.8 M$_\\odot$ of molecular gas with $\\sim$0.5 L$_\\odot$ of mechanical luminosity. HCO$^+$ lines show the kinematic signature of infall motions as well as bulk rotation. The source is most likely a Class 0 protostellar object not yet visible at near-IR wavelengths. With the combination of spatial and spectral resolution in our data set, we are able to disentangle the effects of infall, rotation and outflow towards this young object. ", "introduction": "R Corona Australis (R CrA) is one of the most nearby active star forming regions, at $~\\sim$130pc \\citep{mar81}. This close proximity allows for high spatial resolution observations, even with single dish millimeter wave and submillimeter wave telescopes. The cloud is home to many Herbig Ae/Be and T Tauri stars that have been studied extensively in the near infrared \\citep{wil97}. Recent observations in millimeter wave continuum \\citep{hen94} and molecular lines (\\cite{har93}, \\cite{and97-1}, \\cite{and97-2}), have shown the existence of active, embedded star formation not directly associated with R~CrA, but most likely associated with the Class I source IRS7 \\citep{tay84}. IRS7 is the most deeply embedded near infrared point source in the immediate area. R~CrA was first studied in CO by Loren \\citep{lor79}. His fairly low spatial resolution maps (2.4\\arcmin) showed large amounts of CO peaking at the general location of R~CrA, with high velocity wing emission throughout the map. The wing emission was later interpreted as a molecular outflow by \\cite{lev88}. Levreault's higher spectral and spatial resolution data revealed a large bipolar outflow with an extent of over 10\\arcmin. He also suggested that two outflows might be responsible for the observed morphology, and that the embedded source IRS7 might be responsible for the outflow, not R~CrA itself. In 1985, Wilking et al. used the Kuiper Airborne Observatory to map R~CrA in 100~\\micron~ and 50~\\micron~ continuum emission \\citep{wil85}. The spatial resolution was not adequate to determine if R~CrA, IRS7 or some other unidentified source was driving the continuum and molecular line emission. Since R~CrA and IRS7 are both thought to be high mass sources visible in the near-IR, there is a possibility the driving source is neither of these objects. At both 50 and 100~\\micron~, the emission was found to be extended, with a flux density of 570 Jy/beam at 100~\\micron~. These FIR continuum observations were supplemented by Henning \\citep{hen94}. This 23\\arcsec ~ resolution map spatially resolved the region, and showed that the peak of the FIR emission is not associated with R~CrA, but is much closer to the infrared source IRS7. In 1993, Harju et al. performed a large area (40\\arcmin$\\times$10\\arcmin) survey of the entire R~CrA molecular cloud in C$^{18}$O \\citep{har93}. They estimated that the entire cloud complex contains more than 120 M$_\\sun$ of molecular gas. The densest part of this core, centered around IRS7 and R~CrA contains $\\sim$60 M$_\\sun$ of molecular gas. Their kinematic analysis showed evidence for rotation around IRS7, and also clearly identified outflow lobes in the area. Follow-up work by Anderson \\citep{and97-1} identified a large molecular disk around IRS7, which shows the kinematic signature of rotation. They mapped the region around R~CrA in multiple transitions of HCO$^+$ to trace the dense gas. In HCO$^+$, the outflow lobes have an extent of 4\\arcmin, and show no evidence for the multiple outflows suggested by Levreault. In order to probe the disk dynamics further, the central 2\\arcmin$\\times$2\\arcmin\\ region of R~CrA was mapped in HCO$^+$(3-2). The data show an elongated structure, with three peaks to the northwest, southeast and south of IRS7. A slice in centroid velocity through IRS7 at a position angle of 45 degrees shows signs of rotation. Anderson then followed up their previous work with HCO$^+$(4-3) mapping of the central 2\\arcmin$\\times$2\\arcmin ~ around IRS7, providing improved spatial resolution and kinematic information \\citep{and97-2}. The integrated intensity map shows a structure elongated in the N-S direction, unresolved in the E-W direction, peaking on IRS7. The multi-lobed structure visible in HCO$^+$(3-2) is not present, which they explain as the effect of self-absorption on the line profiles. Velocity centroid maps show a strong velocity gradient near IRS7, and a slice orthogonal to the gradient at a position angle of about 45 degrees revealed the characteristic signature of rotational motion. They fit this rotation curve with a disk+central point source model and derived a central mass of 0.8 M$_\\sun$. This result points towards the conclusion that the mm-wave source is a deeply embedded protostar of fairly low mass. ", "conclusions": "\\subsection{A Possible Class 0 Source Driving the FIR Emission} In most previous research, it is assumed that the molecular line and FIR emission in the region is directly associated with R~CrA. However, work by Harju and Anderson suggested the embedded Class I source IRS7 as the most likely driving source for this emission (\\cite{har93}, \\cite{and97-1}, \\cite{and97-2}). The possibility exists that the peak of the mm-wave and sub-mm wave emission, and the driving source for the molecular outflow(s) in the region might not be IRS7, but a more deeply embedded source not detected in infrared surveys to date. The results presented here suggest that the source of the molecular line and cold dust emission may be a deeply embedded Class 0 source within $\\sim$10\\arcsec\\ of IRS7, with a mass of about 0.5 M$_\\sun$. Recent work by \\cite{chi03} concludes that IRS7 is not the driving source behind the mm-wave emission in the region. Their 1.3mm continuum map of the region identifies 25 dust emission peaks. The peak near IRS7, called MMS13, is conincident with neither the thermal IR source associated with IRS7 from \\cite{wil97} nor the VLA sources from \\cite{bro87}, but is located $\\sim$15\\arcsec\\ south. They conclude that MMS13 is most likely a deeply embedded Class 0 source, although their continuum flux suggests a source mass of 5 M$_\\odot$ for 20K dust. Our data also seem to suggest the driving source is south of IRS7, but our spatial resolution is not adequate to make a firm determination. Recently, 50 GHz continuum interferometric imaging of the region was performed by \\cite{cho04}. They find an elongated structure with a position angle of $\\sim$120 degrees located $\\sim$5\\arcsec\\ to the north of IRS7. The elongated structure has a SED consistent with dust emission, while the emission peak at the position of IRS7 has an SED suggesting free-free or non-thermal emission. Millimeter wave interferometric observations toward IRS 7 in one or more molecular lines (e.g. CO, HCO$^+$ and CS) are needed to identify the location of the driving source. Other evidence exists that suggests the source of the emission is a Class 0 object. The mass of the molecular core can be estimated three ways: 1) a kinematic measure (Figure 4) of the total mass of the protostar and surrounding gas. 2) the FIR continuum flux and 3) the column density of HCO$^+$. All three of these measurements have considerable uncertainty, but taken together point towards a deeply embedded source of about 0.5 M$_\\sun$ with a gaseous envelope of similar mass. The kinematic measurement from the centroid velocity of HCO$^+$(4-3) shows a velocity gradient strongly suggestive of rotation. If this velocity gradient is due to rotation, then a lower limit to the enclosed mass can be derived. In the 0.030 pc diameter volume about the center of motion, we derive a gravitational mass of 0.75 M$_\\sun$. This assumes all the mass is concentrated at the center. We also estimated the gas mass in the central 22\\arcsec\\ (about half the distance between velocity peaks in the velocity gradient) via continuum observations. Assuming the dust emission arises from 0.1 \\micron\\ radius silicate dust grains with a dust emissivity $\\propto \\lambda^{-2}$ and the dust to gas ratio is $\\sim$100, we compute a total gas mass for the core of 0.6 M$_\\odot$. This determination has considerable uncertainy due to the use of IRAS data in the determination of the SED, but is consistent with the masses derived using other methods. The measurement of the $H_2$ mass via HCO$^+$ also gives a similar answer. We used H$^{12}$CO$^+$ and H$^{13}$CO$^+$ measurements with a 22\\arcsec ~ beamsize to estimate the column density of $H_2$. Both the varying abundance of HCO$^+$ to $H_2$ and the effects of self absorption are uncertain. The presence of self-absorption will lead to overestimates of the HCO$^+$ column density. Assuming a HCO$^+$/H$_2$ abundance ratio of 10$^{-9}$ and neglecting the effects of self absorption, we estimate a gas mass of 0.5 M$_\\sun$. While uncertain, both these measures point to a total envelope mass on the order of the mass of the central source. In addition, the somewhat small FIR luminosity of 21 L$_\\odot$ is similar to low mass protostars of about 0.5 M$_\\sun$ \\citep{wal90-2}. We suggest that the driving source for the outflow is a previously unidentified Class 0 protostar located $\\sim$ 10\\arcsec\\ from IRS~7 at the peak of the mm-wave continuum emission, referred to by \\cite{chi03} as MMS13. \\subsection{Infall} Spectral lines of Class 0 protostellar objects sometimes show the classic spectral signature of infall \\citep{wal86}. As seen in Figure 6, HCO$^+$(4-3) spectra near the center position of the map show an enhanced blue peak, characteristic of this infall signature. In addition, the centroid velocity map in Figure 4 shows the ``Blue Bulge'' signature of infall motion \\citep{wal94}. North and south of the central position, rotation begins to dominate the kinematics. At about 30\\arcsec ~ east and west of the source the outflow dominates the appearance of the line profiles in the HCO$^+$ map. In CO maps, the outflow completely masks all signs of rotation and infall. Only in dense gas tracers (e.g. HCO$^+$) do all the intertwined motions reveal themselves near the source. Convolving spectra to larger beamsizes, we continue to see a line shape suggestive of infall. In Figure 7, the HCO$^+$ spectrum has been convolved in half beam increments from 22\\arcsec ~ to 77\\arcsec, which is essentially the entire HCO$^+$(4-3) map. The infall lineshape is still clearly visible in all the spectra, suggesting that infall is occurring over the entire central region around the protostellar source. This points toward an infall region of at least 6000 AU in radius. Using our measurements of FIR luminosity and temperature, we can estimate the mass infall rate for this object. If all the FIR luminosity is generated through accretion, then \\begin{equation} L_{FIR}=\\frac{G M \\dot{M}}{r} \\end{equation} \\noindent where {\\it r} is the radius of the protostar and {\\it M} is the mass of the protostar. We assume $2 \\times 10^{11}$ cm for the radius, about 3 R$_\\odot$ \\citep{sta80}, and 0.5 M$_\\sun$ for the mass of the central object, given our mass estimates of the source+envelope and envelope masses. This leads to a mass accretion rate of $4 \\times 10^{-6}$ M$_\\sun$/yr, uncertain to within a factor of a few. A mass infall rate can also be computer for the case of self-similar collapse. Assuming a soundspeed due to only thermal pressure, and ignoring magnetic and turbulent support, the mass accretion rate is given by \\citep{shu77}: \\begin{equation} \\dot{M}=\\frac{0.975~a^3}{G} \\end{equation} \\noindent where, \\begin{equation} a=\\sqrt{\\frac{k T}{m_{H_2}}} \\end{equation} \\noindent Substituting the gas temperature derived from the observed spectral energy distribution ($\\sim$36K), we find a mass accretion rate of $1 \\times 10^{-5}$ M$_\\sun$/yr. Both these estimates are consistent with a low mass, Class 0 source. \\subsection{Atomic Carbon Distribution} In the canonical picture of carbon in a molecular cloud, an \"onion-skin\" model is used where the carbon is ionized in the outer layer of the cloud as [CII]. Further in the cloud where the carbon is more shielded, it is in atomic form as [CI]. Deeper in the cloud, the carbon forms CO. In some cases, where the cloud material is very clumpy, this same picture holds, but now each clump acts as a small cloud, causing [CII] and [CI] to appear well mixed throughout the cloud \\citep{stu88}. We see evidence for both these effects in our carbon observations. To reveal enhancement of [CI] relative to C$^{18}$O, we calculated the column density of both species, assuming they are both optically thin. We estimated the column density of the atomic carbon following \\cite{wal93-2}: \\begin{equation} N_C=\\frac{N_1}{3} \\left ( e^{\\frac{23.6}{T_{ex}}}+3+5 e^{\\frac{-38.8}{T_{ex}}} \\right ) \\end{equation} \\noindent where \\begin{equation} N_1=5.94 \\times 10^{15} \\int T_{mb}~dV~cm^{-2} \\end{equation} \\noindent We then plotted the ratio of [CI]/C$^{18}$O column density normalized to the average over the map to look for enhancement/depletion of atomic carbon relative to CO, considering only locations were both the [CI] and CO integrated intensities were larger than the 2$\\sigma$ noise level. The results are shown in Figure 8. As expected from the canonical picture, [CI] is enhanced at the edges of the cloud, and depleted towards the center. We have plotted the contours over an optical image of the region from the Palomar Sky Survey. The regions of low relative [CI] abundance correlate with the regions of high optical extinction, and the region with strong FIR continuum emission. In the outflow map of Figure 5, we see that carbon is observed throughout the red wing of the outflow. Only a small amount of carbon is visible in the blue wing. While the [CI] is not as ubiquitous as $^{13}$CO in the outflow, it is still visible, and does not appear to be in a shell configuration. This leads us to believe that while atomic carbon is enhanced at the edge of the cloud, some atomic carbon is still mixed throughout the cloud, pointing to clumpy structure even in the outflow. \\subsection{System Configuration} The kinematic signs of rotation, infall and outflow, combined with the distribution of CO, HCO$^+$ and submillimeter continuum emission all combine to produce a picture of star formation in the R~CrA molecular core. Due to the proximity of R~CrA, the kinematic signatures of infall, rotation and outflow can be tentatively identified and disentangled. A proposed configuration for the MMS13 core is shown in Figure 9, overlaid on the centroid velocity plot of HCO$^+$(4-3). We believe a highly embedded Class 0 protostellar source is the driver for the FIR emission in the area. Kinematic and morphological evidence point to a location within $\\sim$10\\arcsec\\ of the Class I source IRS7. This source, located at the peak of the mm-wave continuum emission, was tenatively identified by \\cite{chi03} as MMS13. This source is surrounded by a rotating molecular disk, first observed by Anderson et al. (\\cite{and97-1}, \\cite{and97-2}). Our HCO$^+$(4-3) measurements point to a similar conclusion, with an enclosed mass of 0.75 M$_\\sun$ $\\pm$ 0.15 M$_\\sun$. Mass estimates based on derived gas column density, 870 \\micron\\ continuum measurements and dynamical arguments suggest a central protostar of $\\sim$0.5 M$_\\sun$, with an envelope of similar mass. We postulate the disk major axis is at a position angle of about 60 degrees. This position angle gives the cleanest rotational signature in the velocity gradient, and is consistent with the emission and the orientation of the molecular outflow. This molecular disk has a major axis of $\\sim$2500 AU, and is unresolved in the orthogonal direction in all our observations. In the central 20\\arcsec-30\\arcsec ~ of the HCO$^+$(4-3) map, the line profiles show a strong infall line shape. This shape persists when the spectra are convolved to beamsizes as large as the entire HCO$^+$ map. In addition, the central region of the centroid velocity plot shows the \"Blue Bulge\" signature of infall motion. This infall region is at least 6000 AU in radius, surrounding the molecular disk. At distances larger that 6000 AU, the kinematics of the molecular outflow begin to dominate the line profiles. At lower spatial resolution, and with more sensitivity to the lower density gas, CO maps are virtually all dominated by the outflow and the ambient cloud material." }, "0405/hep-th0405199_arXiv.txt": { "abstract": "s{ A field theory is studied where the consistency condition of equations of motion dictates strong correlation between states of \"primordial\" fermion fields and local value of the dark energy. In regime of the fermion densities typical for normal particle physics, the primordial fermions split into three families identified with regular fermions. When fermion energy density is comparable with dark energy density, the theory allows transition to new type of states. The possibility of such Cosmo-Low Energy Physics (CLEP) states is demonstrated in a model of FRW universe filled with homogeneous scalar field and uniformly distributed nonrelativistic neutrinos. Neutrinos in CLEP state are drawn into cosmological expansion by means of dynamically changing their own parameters. One of the features of the fermions in CLEP state is that in the late time universe their masses increase as $a^{3/2}$ ($a=a(t)$ is the scale factor). The energy density of the cold dark matter consisting of neutrinos in CLEP state scales as a sort of dark energy; this cold dark matter possesses negative pressure and for the late time universe its equation of state approaches that of the cosmological constant. The total energy density of such universe is less than it would be in the universe free of fermionic matter at all.} ", "introduction": " ", "conclusions": "" }, "0405/hep-th0405016_arXiv.txt": { "abstract": "We review the recent attempts of unifying inflation with quintessence. It appears natural to join the two ends in the framework of brane world cosmology. The models of quintessential inflation belong to the class of {\\it non-oscillatory} models for which the mechanism of conventional reheating does not work. Reheating through gravitational particle production is inefficient and leads to the excessive production of relic gravity waves which results in the violation of nucleosynthesis constraint. The mechanism of {\\it instant preheating} is quite efficient and is suitable for brane world quintessential inflation. The model is shown to be free from the problem of excessive production of gravity waves. The prospects of Gauss-Bonnet brane world inflation are also briefly indicated. ", "introduction": "Universe seems to exhibit an interesting symmetry with regard to accelerated expansion. It has gone under inflation at early epochs and is believed to be accelerating at present. The inflationary paradigm was originally introduced to address the initial value problems of the standard hot big bang model. Only later it became clear that the scenario could provide important clues for the origin of structure in the universe. The recent measurement of the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{WMAP1,WMAP2} in the Cosmic Microwave Background (CMB) made it clear that (i) the current state of the universe is very close to a critical density and that (ii) primordial density perturbations that seeded large-scale structure in the universe are nearly scale-invariant and Gaussian, which are consistent with the inflationary paradigm. Inflation is often implemented with a single or multiple scalar-field models\\cite{LR} (also see the excellent review on inflation by Shinji Tsujikawa\\cite{shinji}). In most of these models, the scalar field undergoes a slow-roll period allowing an accelerated expansion of the universe. After drawing the required amount of inflation, the inflaton enters the regime of quasi-periodic oscillation where it quickly oscillates and decays into particles leading to (p)reheating.\\par As for the current accelerating of universe, it is supported by observations of high redshift type Ia supernovae treated as standardized candles and, more indirectly, by observations of the cosmic microwave background and galaxy clustering. Within the framework of general relativity, cosmic acceleration should be sourced by an energy-momentum tensor which has a large negative pressure (dark energy)\\cite{phiindustry}. Therefore, the standard model should, in order to comply with the logical consistency and observation, be sandwiched between inflation at early epochs and quintessence at late times. It is natural to ask whether one can build a model with scalar fields to join the two ends without disturbing the thermal history of universe. Attempts have been made to unify both these concepts using models with a single scalar field \\cite{unifiedmodels}.In these models, the scalar field exhibits the properties of tracker field. As a result it goes into hiding after the commencement of radiation domination; it emerges from the shadow only at late times to account for the observed accelerated expansion of universe. These models belong to the category of {\\it non oscillating} models in which the standard reheating mechanism does not work. In this case, one can employ an alternative mechanism of reheating via quantum-mechanical particle production in time varying gravitational field at the end of inflation \\cite{ford}. However, then the inflaton energy density should red-shift faster than that of the produced particles so that radiation domination could commence. And this requires a steep field potential, which of course, cannot support inflation in the standard FRW cosmology. This is precisely where the brane\\cite{randall,h} assisted inflation comes to the rescue.\\par The presence of the quadratic density term (high energy corrections) in the Friedman equation on the brane changes the expansion dynamics at early epochs \\cite{cline}(see Ref\\cite{roy rev} for details on the dynamics of brane worlds) Consequently, the field experiences greater damping and rolls down its potential slower than it would during the conventional inflation. Thus, inflation in the brane world scenario can successfully occur for very steep potentials\\cite{basset,liddle}. The model of quintessential inflation based upon reheating via gravitational particle production is faced with difficulties associated with excessive production of gravity waves. Indeed the reheating mechanism based upon this process is extremely inefficient. The energy density of so produced radiation is typically one part in $10^{16}$\\cite{liddle} to the scalar-field energy density at the end of inflation. As a result, these models have prolonged kinetic regime during which the amplitude of primordial gravity waves enhances and violates the nucleosynthesis constraint\\cite{vst}(see also \\cite{star79}). Hence, it is necessary to look for alternative mechanisms more efficient than the gravitational particle production to address the problem. A proposal of reheating with Born-Infeld matter was made in Ref\\cite{bisn}(see also Ref\\cite{bi1,bi2} on the related theme). It was shown that reheating is quite efficient and the model does not require any additional fine tuning of parameters\\cite{bisn}. However, the model works under several assumptions which are not easy to justify.\\par The problems associated with reheating mechanisms discussed above can be circumvented if one invokes an alternative method of reheating, namely `instant preheating' proposed by Felder, Kofman and Linde \\cite{FKL} (see also Ref\\cite{shtanov} on the related theme. For other approaches to reheating in quintessential inflation see \\cite{curvaton}). This mechanism is quite efficient and robust, and is well suited to non-oscillating models. It describes a new method of realizing quintessential inflation on the brane in which inflation is followed by `instant preheating'. The larger reheating temperature in this model results in a smaller amplitude of relic gravity waves which is consistent with the nucleosynthesis bounds\\cite{samiv}. However, the recent measurement of CMB anisotropies by WMAP places fairly strong constraints on inflationary models \\cite{spergel03,tegmark03}. It seems that the steep brane world inflation is on the the verge of being ruled out by the observations\\cite{suji04}. Steep inflation in a Gauss-Bonnet braneworld may appear to be in better agreement with observations than inflation in a RS scenario \\cite{lidsey}. ", "conclusions": "In this paper we have reviewed the recent work on unification of inflation with quintessence in the frame work of brane worlds. These models belong to the class of {\\it non-oscillatory} models in which the underlying alternative reheating mechanism plays a crucial role. The popular reheating alternative via quantum mechanical production of particle during inflation leads to an unacceptable relic gravity wave background which violates the nucleo-synthesis constraint at the commencement of radiative regime. We have mentioned other alternatives to conventional (p)reheating and have shown that 'instant preheating' discovered by Felder, Linde and Kofman is superior and best suited to brane world models of quintessential inflation. The recent measurement of CMB anisotropies by WMAP, appears to heavily constraint these models. The steep brane world inflation seems to be excluded by observation in RS scenario\\cite{suji04}. As shown in Ref.\\cite{ssr}, the inclusion of GB term in the bulk effects the constraints on the inflationary potentials and can rescue the steep exponential potential allowing it to be compatible with observations for a range of energy scales. The GB term leads to an increase of the spectral index $n_S$ and decrease of tensor to scalar ratio of perturbations $R$ in the intermediate region between RS and GB. As seen from Fig.\\ref{exponential.ps}, there is an intermediate region where the steep inflation driven by exponential potential lies within $2 \\sigma$ contour for ${\\cal N}=70$ . Thus, the steep inflation in a Gauss-Bonnet braneworld appears to be in agreement with observations." }, "0405/astro-ph0405084_arXiv.txt": { "abstract": "A five-dimensional Ricci-flat cosmological solution is studied by assuming that the induced 4D matter contains two components: the usual fluid for dark matter as well as baryons and a scalar field with an exponential potential for dark energy. With use of the phase-plane analysis it is shown that there exist two late-time attractors one of which corresponds to a universe dominated by the scalar field alone and the other is a scaling solution in which the energy density of the scalar field remains proportional to that of the dark matter. It is furthermore shown that for this 5D scaling solution the universe expands with the same rate as in the 4D FRW models and not relies on which 4D hypersurface the universe is located in the 5D manifold. ", "introduction": "Scalar fields play a central role in modern cosmology in driving inflation of the early universe and describing dark energy of the present epoch.$% ^{1,2} $ Observations predict that our universe contains roughly one-third of dark matter and baryons and two- thirds of dark energy.$^{3}$ Within the standard Friedmann-Robertson-Walker (FRW) models it was shown that there exist scaling solutions that are the unique late-time attractors whenever they exist.$^{4,5,6}$ Here, in this paper, we wish to look for scaling solutions in higher-dimensional cosmological models. In Kaluza-Klein theories as well as in brane world scenarios, our 4D universe is believed to be embedded in a higher-dimensional manifold. One of these models is the Ricci-flat 5D cosmological solutions presented by Liu and Wesson.$^{7}$ This model is 5D Ricci-flat, implying that it is empty viewed from 5D. However, as is known from the induced matter theory,$^{8,9}$ 4D Einstein equations with matter could be recovered from 5D equations in apparent vacuum. This approach is guaranteed by Campbell's theorem that any solution of the Einstein equations in N-dimensions can be locally embedded in a Ricci-flat manifold of (N+1)-dimensions.$^{10}$ In section II, we suppose the 4D induced matter be composed of a perfect fluid and a scalar field. In section III, we use phase-plane analysis to study the evolutions of the model. In section IV we study the scaling solution. Section V is a short discussion. ", "conclusions": " (1). The scaling solution (\\ref{p}), as a late-time attractor of the evolution of the universe, is the same in both the 4D FRW models and the 5D induced matter theory. (2). In the 5D model (\\ref{t})-(\\ref{w}), the expansion rate $A$ and $H$ of (% \\ref{t}) seems not rely on the value of $y$. In other words, our 4D universe, no matter it is situated on which hypersurface of $y$ in the 5D manifold, looks similar and expands with almost the same rate, in late times. (3). The constant $K$ in (\\ref{u}) represents the curvature of the 5D manifold (see Eq.(\\ref{b})). It is reasonable to believe that $K=0,\\pm 1$ may correspond to three different topologies for the 5D manifold (\\ref% {5Dmatric}). However, from (\\ref{t})-(\\ref{w}), we find that $K$ does not affect $A$ and $H$ so much in late-times. (4). In Eq. (\\ref{t}), if $\\gamma _{m}=1/3$, we have $B\\approx 1$ and $% A\\approx t^{2}$, and our model approaches to the Milne model which has no event or particle horizon, and so it is of particular interest. It was shown$^{7,12}$ that in the early-time of the universe the 5D solution (\\ref{5Dmatric}) deviates from the standard FRW models greatly. For instance it may have a big bounce rather than a big bang and before the bounce the universe is contracting. Meanwhile, the value of the 5D curvature $K$ is sensitive for having or not having a bounce. However, in the late time, we have $B\\longrightarrow 1$ and the coordinate time $t$ gives back to the cosmic time $\\tau $,\\ so the bounce model approaches to the FRW model. As a late-time attractor, the scaling solution seems not show up noticeable deviations from the standard FRW models. Further study along this line is needed to distinguish the two theories." }, "0405/astro-ph0405567_arXiv.txt": { "abstract": "We present the results of an observing program with the SCUBA bolometer array to measure the submillimetre (submm) dust continuum emission of \\nSMMobsz\\ distant ($z > 1$) radio galaxies. We detected submm emission in \\nSMMdetzR\\ galaxies with S/N $> 3$, including 9 detections at $z > 3$. When added to previous published results these data almost triple the number of radio galaxies with $z > 3$ detected in the submm and yield a sample of \\nSMMobstotAR\\ observed radio galaxies over the redshift range $z$ = 1--5. We find that the range in rest-frame far-infrared luminosities is about a factor of 10. We have investigated the origin of this dispersion, correlating the luminosities with radio source power, size, spectral index, $K$-band magnitude and \\lya\\ luminosity. No strong correlations are apparent in the combined data set. We confirm and strengthen the result from previous submm observations of radio galaxies that the detection rate is a strong function of redshift. We compare the redshift dependence of the submm properties of radio galaxies with those of quasars and find that for both classes of objects the observed submm flux density increases with redshift to $z \\approx 4$, beyond which, for the galaxies, we find tentative evidence for a decline. We find evidence for an anti-correlation between submm luminosity and UV polarisation fraction, for a subsample of 13 radio galaxies, indicating that starbursts are the dominant source of heating for dust in radio galaxies. ", "introduction": "There is strong evidence that powerful high redshift radio galaxies (\\hzrgs; $z > 2$) are the progenitors of the brightest cluster ellipticals seen today. \\hzrgs\\ are the infrared brightest and presumably the most massive galaxies at any epoch \\citep{DeBreuck02aj} and host actively-accreting super massive black holes with masses of order $10^{9}\\,\\Msun$ \\citep{Lacy01apj,Dunlop03mnras}. Therefore, they are key objects for studying the formation and evolution of massive galaxies and super-massive black holes. \\hzrgs\\ are likely to be in an important phase of their formation process for several reasons: They have large reservoirs of gas from which they could be forming, as shown by spectacular ($>$ 100 kpc) luminous \\lya\\ haloes \\citep[\\eg][]{McCarthy93araa,vanOjik96aa,Reuland03apj} and widespread \\hi\\ absorption features in the \\lya\\ profiles \\citep{vanOjik97aa}. Their rest-frame UV morphologies are characterized by clumpy structures, similar to the Lyman-break galaxies at $z \\sim 3$ , that will merge with the central galaxy on dynamical time-scales of $10^8$\\,yrs \\citep{Pentericci98apj,Pentericci99aa}. In the case of 4C~41.17 there is direct evidence for massive star formation (up to $\\sim 1500\\,\\Msunpyr$ after correction for extinction) based on stellar absorption--lines \\citep{Dey97apj}. Finally, mm-interferometry studies of CO line and continuum emission for three $z > 3$ HzRGs have shown that the star formation occurs galaxy wide over distances up to 30\\,kpc \\citep{Papadopoulos00apj,DeBreuck03aa} and there is even evidence for star formation on scales of 250\\kpc\\ \\citep{Stevens03nat}. Together this suggests that we are observing not merely scaled up versions of local ultraluminous infrared galaxies (ULIRGs) where the bursts are confined to the inner few kpc, but wide-spread starbursts within which the galaxies are forming the bulk of their eventual stellar populations. \\hzrgs\\ are an important sample for studying the star formation history of the universe because their selection is based on long wavelength radio emission whose propagation is not affected by the presence of dust. Dust is expected to play a significant role in star forming regions, absorbing UV/optical radiation from the starburst and reradiating it at far-infrared (FIR) wavelengths \\citep{SandersMirabel96araa}. Optical searches for distant galaxies \\citep[\\eg\\ using the Lyman-break technique;][]{Steidel96apj, Steidel99apj,Ouchi01apj} are thus likely to be biased against dusty objects. Finding distant star forming galaxies through submillimetre (submm; rest-frame FIR) emission \\citep[\\eg\\ ][]{Hughes98nat,Bertoldi02confproc,Chapman02amnras,Cowie02aj,Scott02mnras,Smail02mnras,Webb03aapj,Eales03mnras} selects only the most obscured sources. So far there has been little overlap between the optical and submm selected star forming sources \\citep[selection on very red near-IR colours may prove more fruitful; \\eg\\ ][]{Frayer04aj}. It remains unclear whether they are members of a continuous population \\citep[\\eg ][]{AdelbergerSteidel00apj,Webb03bapj} and arguments have been made that either could dominate the star formation density at high redshift \\citep{Blain99mnras,AdelbergerSteidel00apj}. Since selection at radio wavelengths circumvents the aforementioned selection biases it could help determine the relative contributions of obscured and unobscured star formation to the star formation history of the universe. \\citet[][hereafter A01]{Archibald01mnras} have conducted the first systematic submm survey to study the star formation history of radio galaxies over a redshift interval of $0.7 < z < 4.4$. In their sample of 47 galaxies, they found evidence for a considerable range in FIR luminosities, a substantial increase in 850\\,\\micron\\ detection rate with redshift and that the average 850\\,\\micron\\ luminosity rises at a rate $(1 + z )^{3-4}$ out to $z \\simeq 4$. These results prompt the following questions: Is the dispersion in FIR luminosities due to differences in their star formation rates or dust contents? Does the strong increase with redshift reflect an increase in star formation rates or could it be related to changing dust properties? Does the FIR luminosity keep on rising with redshift or does it level off and is there a redshift cut-off? Are the inferred star formation rates comparable to those derived from the optical/UV? Do the submm properties of quasars (QSOs) and radio galaxies show similar trends or do the two classes of objects evolve differently? The submm findings from A01 were based on a limited number of detections at high redshift ($z > 3$). To put these results on a statistically firmer footing and search for possible correlations with other galaxy parameters more submm observations were required. Here we present such observations of all $z > 3$ HzRGs known at the beginning of 2001 \\citep[\\eg][]{DeBreuck01aj} which had not been observed in the submm. Adding these to the survey of A01 almost triples the number of detections at high redshift, creating a sample which is statistically significant over the full redshift range $z =$ 1 -- 5. The structure of this paper is a follows: the sample selection, observations and data analysis are described in Section 2. Results and notes on some individual sources are presented in Section 3. Various correlations with submm properties of \\hzrgs\\ are investigated in Section 4 and described in detail in Section 5. Section 6 presents a comparison between \\hzrgs\\ and QSOs. We discuss and summarize our conclusions in Section 7. Throughout this paper, we adopt a flat universe with $\\OmM = 0.3$, $\\OmL = 0.7$, and $H_{0} = 65 \\kmps\\,\\rm Mpc^{-1}$. Using this cosmology the look-back time at $z \\sim 2.5$ (the median redshift of our sample) is 11.7\\,\\ph\\,Gyr and a galaxy at such a redshift must be less than 2.8\\,\\ph\\, Gyr old. ", "conclusions": "We have presented SCUBA observations of 24 radio galaxies and compared those with earlier results of 47 radio galaxies from the survey by A01. We confirm that \\hzrgs\\ are massive forming galaxies, forming stars up to rates of a few thousand \\Msunpyr\\ and that there is no strong evidence for a correlation with radio power. Further evidence for a predominantly starburst nature of the far-IR emission comes from the striking anti-correlation between submm flux density and UV polarisation (Fig. \\ref{UVpolS850}). In agreement with A01 we find that submm detection rate appears to be primarily a function of redshift. If this is interpreted as being due to a change in the intrinsic far-IR luminosity, it would be consistent with a scenario in which the bulk of the stellar population of radio galaxies forms rapidly around redshifts of $z = 3-5$ after which they are more passively evolving \\citep[c.f. ][]{Best98mnras}. We also find that the median redshift of the \\hzrgs\\ with SCUBA detections ($z = 3.1$) is higher than the median redshift of the submm population \\citep[$z = 2.4$;][]{Chapman03nat}. In the current picture of hierarchical galaxy formation, this could be interpreted as that \\hzrgs\\ are more massive galaxies, which are then thought to begin their collapse at earlier cosmic times and evolve faster and finish the bulk of their formation process earlier. Alternatively, it could indicate that higher redshift submm sources are being missed due to the requirement of a radio counterpart prior to spectroscopic follow-up. \\hzrgs\\ have accurately determined redshifts and host identifications and are thought to be the most massive galaxies at any epoch \\citep{DeBreuck02aj}. Therefore, \\hzrgs\\ are a key population for studies of galaxy formation in the early universe, allowing detailed follow-up mm-interferometry observations to study their dust and gas content. Currently, they offer the best way to obtain reliable dynamic masses for a significant number of massive high redshift galaxies. These \\hzrgs\\ would be especially suited to constrain any evolution in galaxy mass with redshift, study changes in evolutionary status, gas mass, and the starburst-AGN connection. If all of these turn out to have masses larger than $10^{11}$\\Msun, then this could have important consequences (depending on rather uncertain correction factors for the fraction of similar galaxies for which the black hole is dormant) for our understanding of galaxy formation, because only few such massive galaxies are expected at such high redshifts \\citep{Genzel03apj}." }, "0405/astro-ph0405421_arXiv.txt": { "abstract": "We present results from a deep photometric study of the rich galaxy cluster Abell 2218 ($z=0.18$) based on archival HST WFPC2 F606W images. These have been used to derive the luminosity function to extremely faint limits ($M_{\\rm F606W} \\approx -13.2$ mag, $\\mu_{0}\\approx 24.7$\\,mag\\,arcsec$^{-2}$) over a wide field of view ($1.3\\; h^{-2}\\;$Mpc$^{2}$). We find the faint-end slope of the luminosity function to vary with environment within the cluster, going from $\\alpha = -1.23\\pm 0.13$ within the projected central core of the cluster ($100 < r < 300\\,h^{-1}$\\,kpc) to $\\alpha = -1.49\\pm 0.06$ outside this radius ($300 < r < 750\\,h^{-1}$\\,kpc). We infer that the core is 'dwarf depleted', and further quantify this by studying the ratio of `dwarf' to `giant' galaxies and its dependency as a function of cluster-centric radius and local galaxy density. We find that this ratio varies strongly with both quantities, and that the dwarf galaxy population in A2218 has a more extended distribution than the giant galaxy population. ", "introduction": "The galaxy luminosity function (LF) -- the number density of galaxies per unit luminosity interval -- is a fundamental descriptor of the galaxy population and, as such, contains important information on the formation and evolution of galaxies (Bingelli, Sandage \\& Tammann 1988, Benson et al. 2003). At low redshift, the LF will contain the combined imprints of the galaxy initial mass function (Press \\& Schechter 1974, Schechter 1976), together with the effects of any subsequent evolutionary processes which modified this distribution (e.g., hierarchical merging; White \\& Rees 1978). Determinations of the LF in different environments and at a variety of redshifts provide the only direct hope of disentangling these environmental and evolutionary effects. In this context, rich clusters are fundamental testing grounds, representing the densest environments in which galaxies reside and, within the hierarchical clustering framework, the ultimate examples of where build-up through merging has taken place. Here galaxies at {\\it all} luminosities, both bright and faint, may bear evidence of environmental effects: the `giant' ($L>L^*$) galaxies at the bright end of the LF are the strongest candidates for having been built up through successive mergers and accretion (e.g. Kauffmann, White \\& Guiderdoni 1993). Equally, the sub-luminous ($L0.3$ are few (Nelson et al. 2001; Barrientos \\& Lilly 2003). The present paper addresses this issue and combines a discussion of LF evolution in cluster galaxies with a simultaneous assessment of the colour--magnitude relation. Data acquisition and reduction is discussed in Section 2. Section 3 presents the colour--magnitude relation, colour distribution and the LF computed for individual clusters. Section 3 also discusses the evolution of the above quantities with redshift. The main results are presented and discussed in Section 4, i.e. we demonstrate that the cluster galaxy population is consistent with the presence of at least two components: old systems formed at high redshift that have evolved passively from that epoch, together with a galaxy population displaying more recent star formation. We discuss constraints placed on the evolution of both populations by the current data set. Throughout this paper we assume a Friedmann--Robertson--Walker--Lemaitre cosmological model described by the parameters $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70$ kms$^{-1}$ Mpc$^{-1}$. ", "conclusions": "We have studied a sample of 24 clusters located at redshifts $0.2970.4$ and 6 have $z>0.6$. The majority of the clusters are either X--ray selected or detected, and we are therefore observing gravitationally bound systems. Most of the cluster sample, particularly clusters at redshifts $z<0.6$, possess X--ray luminosities and optical richness values typical of groups or low mass clusters. All clusters in our sample, despite the primary X--ray selection and low X--ray flux/optical richness displayed by the majority of the sample, display a statistical overdensity of galaxies of similar colour (Figure 3), that make them detectable by an almost 3 dimensional search defined by sky position and colour. In fact, all clusters with $R$ and $z'$ photometry, with the exception of XLSSC 007, are colour--detected. However, the present optical identification of XLSSC 007 as the counterpart of the extended X--ray source is uncertain as a result of the large distance between the optical overdensity and X--ray centres. Should the identification of this X--ray source change to that of a $z>1$ cluster, then no X--ray selected cluster presented in this paper is missed by the $R-z'$ technique in the $z<1$ regime. Most of the clusters are identified in X--rays, largely independent of the optical luminosity of the member galaxies. Therefore, the colour detection is non--trivial. The majority of the clusters are optically poor (Abell richness class 0 or lower) consistent with the low computed X--ray luminosities. We have therefore demonstrated that a colour plus spatial overdensity search technique can effectively identify optically poor systems at intermediate to high redshifts (at least those previously identified in X--rays). The emerging picture from the current study is the one of a typical cluster composed of two or more distinct galaxy populations : a relatively old population evolving passively (as measured from the evolution of the color of the red sequence) together with a younger population, ostensibly responsible for the apparent brightening of the characteristic LF magnitudes. The reddest galaxies within each cluster/group evolve in a manner consistent with a model early--type galaxy formed between redshifts $2\\la z_f \\la 5$ (Figure 4). This observation is largely in agreement with previous studies. We note, however, that previous studies employ cluster samples dominated by optically rich systems often observed with heterogeneous instruments. In contrast, the current study consists of an exceptionally uniform cluster sample observed under largely uniform conditions. Previous studies estimate similar values for $z_f$, though assuming different cosmological models, evolutionary models, or both. We note that the formation epoch of cluster galaxies estimated for the current sample would correspond to a higher redshift for the same assumptions adopted in literature. The younger population is detected by studying the LF. The LF of each cluster has been computed in $R-, I-$ and $z'$--bands and is displayed in Figures 5 to 7. A Schechter function provides an acceptable description of the LF shape over the magnitude range extending from $m6''$) that have been observed (see, e.g., Inada et al. 2003). Turner, Ostriker, \\& Gott (1984, hereafter TOG) have considered a simple model in which lensing galaxies, modeled as singular isothermal spheres, are embedded into uniform sheets of background matter. They showed that the effect of the background matter can be important, but also indicated that this effect can be weakened by the presence of low-density regions along the same line of sight. Treating the background matter as a uniform sheet allows an analytical treatment, leading to orders of magnitude estimates, but it is a rather crude representation of the actual matter distribution in a CDM universe in which large-scale structure forms by the growth of Gaussian random fluctuations. In such universe, the rms density fluctuation $\\sigma_{\\rm rms}$ at a particular scale $\\lambda$ is given by \\begin{equation} \\sigma_{\\rm rms}^2(\\lambda) ={1\\over2\\pi^2}\\int_0^\\infty P(k)W^2(k\\lambda)k^2dk\\,, \\end{equation} \\noindent where $P(k)$ is the power spectrum, and $W$ is the top-hat filter function, given by \\begin{equation} W(x)={3\\over x^3}(\\sin x-x\\cos x)\\,. \\end{equation} \\noindent In a CDM universe, small scales collapse before large ones, hence there is at any epoch a maximum scale $\\lambda_{\\max}$ corresponding to the largest structures one expect to find at that epoch in a typical region of the universe. To estimate $\\lambda_{\\max}$, we make the common assumption that the largest structures found in a representative region of the universe correspond to $3\\sigma$ fluctuations. Then, according to the Press-Schechter approximation, the radius of these structures is given by solving \\begin{equation} \\sigma_{\\rm rms}(\\lambda_{\\max})={\\delta_{\\rm crit}\\over3} \\end{equation} \\noindent for $\\lambda_{\\max}$, where $\\delta_{\\rm crit}\\approx1.69$. For a $\\Lambda$CDM universe, we get $\\lambda_{\\max}\\approx23\\,\\rm Mpc$. The typical diameter of the largest objects (voids or clusters) we expect to find in a representative region of the universe is therefore of order $2\\lambda_{\\max}\\approx46\\,\\rm Mpc$. Sources that are gravitationally lensed are normally located at cosmological distances that greatly exceeds that scale. Hence, there will be several overdensities and underdensities along the line of sight to any source, and one could expect a near-cancellation of their effects. TOG acknowledge that fact. Two arguments can be made against this line of reasoning. First, while this near-cancellation is expected for a typical line of sight, it might not occur for an atypical one (this is the argument made by TOG). Second, even though there might be several overdensities and underdensities along the line of sight, their relative contributions to lensing will differ. It is well known that the matter that contributes most to lensing tends to be located about half-way between the source and the observer. To be more quantitative, consider a lens L with projected surface density $\\sigma$, located along the line of sight to a distant source S. We define the convergence $\\kappa$ of the lens as \\begin{equation} \\kappa={\\sigma\\over\\sigma_{\\rm crit}}\\,, \\end{equation} \\noindent where $\\sigma_{\\rm crit}$ is the critical surface density, defined by \\begin{equation} \\sigma_{\\rm crit}={c^2D_S\\over4\\pi GD_LD_{LS}}\\,, \\end{equation} \\noindent and $D_S$, $D_L$, and $D_{LS}$ are the angular diameter distance between observer and source, observer and lens, and source and lens, respectively (SEF, p.~158). The value of $\\kappa$ provides an estimate of the strength of gravitational lensing. We can define a {\\it lensing weight} $w(z)$ using \\begin{equation} w(z)={1\\over\\sigma_{\\rm crit}}={4\\pi GD_LD_{LS}\\over c^2D_S}\\,. \\end{equation} \\noindent The lensing weight gives a measure of the relative contributions to lensing of matter located at various redshifts $z$, all other things being equal. This function is zero at the location of the observer and the source, and peaks at intermediate distances (see Fig.~2 below). The largest contribution to lensing will usually come from matter located in the redshift interval where $w(z)$ is large. The extent of this region is smaller than the whole distance between source and observer. If it is only a few times $\\lambda_{\\max}$, it will contain only a few overdensities and underdensities, and the cancellation of their effect will not be perfect in general. Notice that this effect is partially compensated by the fact that at redshifts $z>0$ where $w(z)$ peaks, the value of $\\lambda_{\\max}$ is smaller than at the present in a CDM universe. There is a clear need for a more realistic approach to this problem than the uniform-sheet approximation of TOG. A recent attempt to quantify the importance of environmental effects on strong lensing by massive galaxies was presented by Holder \\& Schechter (2003). These authors used simulated large-scale structure and galaxy distributions (as we do here), identified the most massive galaxies as being the ones capable of producing strong lensing, and calculated, at the location of these galaxies, the gravitational shear caused by nearby large-scale structures and galaxies. They concluded that the typical shear is of order $10-15\\%$. Our goal in this paper is similar. We want to determine the typical magnitude of environmental effects, and demonstrate that very strong effects are rare. However, we consider an approach to this problem that drastically differs from the one used in previous studies. We use a multiple lens-plane algorithm to simulate the actual gravitational lensing of a large number of distant sources. We performed 100 simulations, each simulation producing the images of 841 sources, for a total of 84,100 images. We extract from these simulations a subsample of images that constitute examples of strong lensing, and we reanalyze this subsample on a case-by-case basis. While the original sample of 84,100 images is unbiased and statistically significant, the subsample of 16 cases is strongly biased, since it includes only cases of strong lensing. Such biased sample is ideal for determining an upper limit to the effects of environment on strong lensing. This work differs from previous studies in our approach to analyzing the effect of the density fluctuations along the line of sight. In particular, instead of focusing on their cumulative effect, we divide these fluctuations in different density components, and determine their relative contributions independently. The most prominent components are clusters of galaxies. We divide these clusters into two components: the galaxies themselves and the diffuse CDM halo in which the galaxies are embedded, and we study the effects of these two components, both individually and in combination. We also consider the effect resulting from the random alignment of galaxies that are physically unassociated. The remainder of this paper is organized as follow: In \\S2, we briefly describe our numerical algorithm. In \\S3, we describe our original set of simulations, and the 16 cases of strong lensing that we have selected for further analysis. In \\S4, we identify, for each case, the particular galaxy (referred to as {\\it The Lens}) that is primarily responsible for producing strong lensing. In \\S5, we investigate the effects of the various density components along the line of sight. In \\S6, we experiment with variations of the basic model, by considering different density profiles for galaxies. Conclusions are presented in \\S7. ", "conclusions": "We have performed a series of ray-tracing experiments using a multiple lens-plane algorithm. We selected 16 cases of strong lensing. By following the trajectory of the beam from the source to the observer, we were able to determine which galaxies along the line of sight were directly hit by the beam. The galaxy that produced the largest value of the convergence $\\kappa$ was identified as The Lens, the galaxy primarily responsible for strong lensing. Our objective was then to study the effects and relative importance of the various components along the line of sight, including (i) The Lens itself, (ii) the galaxies located at the same redshift as The Lens, and possibly associated with it (i.e. in the same cluster), (iii) the background matter located at the same redshift as The Lens, and (iv) the matter located at different redshifts. With the exception of case M, The Lens was always located in the redshift interval $0.3\\leq z\\leq1.1$, where the lensing weight $w(z)=1/\\sigma_{\\rm crit}$ is large (see Fig.~2). With the exception of cases E and M, The Lens was always in a region where the background density is large (see Fig.~6). This was expected; strong lensing is caused by massive elliptical and S0 galaxies, which tend to be located in dense environments according to the morphology-density relation. In 13 cases out of 16, The Lens was among the 5 most massive galaxies within a comoving distance of $5\\,\\rm Mpc$. The effect of the galaxies and background matter associated with The Lens are rather small. The magnifications and image separations vary by $\\sim1\\%$ when the background matter is added, and $\\sim10\\%$ when the galaxies are added, for all cases. In some cases (B, K, and P), a deformation of the image is clearly visible (Fig.~3). Adding the other lens planes produces effects that range from insignificant to spectacular. Insignificant effects occur when The Lens is the only large mass concentration hit by the beam, while spectacular effects occur when several galaxies are hit by the beam, including galaxies as massive or more massive than The Lens. The presence of additional galaxies along the line of sight can create additional images (cases A, M, and O), turn arcs into rings (cases D and N), and cause large increases in magnification, up to factors of several. The effect of the other lens planes on the image separation is less spectacular, but the accumulated effect is still larger than the effect of the galaxies and background matter associated with The Lens. Based on these results, we can summarize the effect of the various components along the line of sight, relative to the effect of The Lens itself, as follows: \\begin{itemize} \\item Background matter located near The Lens: a few percent (difference between squares and triangles in Figs.~4 and 8). \\item Galaxies near The Lens: several percents, up to $10-15\\%$ (cases K and P) (differences between open and filled symbols in Figs.~4 and 8). \\item Galaxies and background matter on other planes: factor of several (difference between filled squares and asterisks in Figs.~4 and 8). \\end{itemize} We conclude that environmental effects usually play a minor role in strong gravitational lensing. The large magnification associated with strong lensing results primarily from a single, massive galaxy (The Lens), or from the random alignment of several physically unassociated galaxies at different distances. It is important to insist that this conclusion is restricted to strong lensing, and cannot be generalized to weak lensing. We showed that environmental effects can be of order 10\\%. If, say, a massive galaxy, acting as a lens, increases the brightness of a distant source by 10\\%, and the nearby galaxies and background matter increase it by an additional 10\\%, the net effect is increased by a factor of 2. Hence, environmental effects can be very important for weak lensing, but are usually not important for strong lensing. We have experimented with various density profiles of galaxies, and reached the same conclusion. In all these experiments, {\\it we have not found one single case for which the nearby background matter or nearby galaxies make any significant difference.} One could debate the statistical significance of our results. We have considered a subsample of 16 images. Had we considered a larger subsample, we might have found a case for which environmental effects are important. For instance, we might find a cluster containing two massive galaxies that happen to be aligned with the source, so that the beam hits both galaxies near their center. In this case, one galaxy would be identified as The Lens, and the effect of the other galaxy would be very important. Such cases must be very rare, though. The fact that a subsample of 16 cases has not turned up one single case for which environmental effects are important, in spite of the fact that our sample only included cases of very strong lensing, suggests that any case for which such effects are important must be atypical. In this, we reach the same conclusion as TOG. Note that this results justifies {\\it a posteriori\\/} our approach of using a biased subsample of 16 cases. If we had found, say, 10 or 11 cases for which environmental effects were very strong, the whole approach would have broken down, but no such cases were found." }, "0405/astro-ph0405329_arXiv.txt": { "abstract": "Motivated by our previous paper, in which we argued for the formation of molecular clouds from large-scale flows in the diffuse galactic interstellar medium, we examine the formation of molecular gas behind shocks in atomic gas using a one-dimensional chemical/dynamical model. In our analysis we place particular emphasis on constraints placed on the dynamical evolution by the chemistry. The most important result of this study is to stress the importance of shielding the molecular gas from the destructive effects of UV radiation. For shock ram pressures comparable to or exceeding typical local interstellar medium pressures, self-shielding controls the formation time of molecular hydrogen but CO formation requires shielding of the interstellar radiation field by dust grains. We find that for typical parameters the molecular hydrogen fractional abundance can become significant well before CO forms. The timescale for (CO) molecular cloud formation is not set by the H$_2$ formation rate on grains, but rather by the timescale for accumulating a sufficient column density or extinction, $A_V \\gtrsim 0.7$. The local ratio of atomic to molecular gas (4:1), coupled with short estimates for the lifetimes of molecular clouds (3-5 Myr), suggests that the timescales for accumulating molecular clouds from atomic material typically must be no longer than about 12-20 Myr. Based on the shielding requirement, this implies that the typical product of pre-shock density and velocity must be $n v \\gtrsim 20\\;\\cc\\ {\\rm km s^{-1}}$. In turn, depending upon the shock velocity, this implies shock ram pressures which are a few times the typical estimated local turbulent gas pressure, and comparable to the total pressures (gas plus magnetic plus cosmic rays). Coupled with the rapid formation of CO once shielding is sufficient, flow-driven formation of molecular clouds in the local interstellar medium can occur sufficiently rapidly to account for observations. We also provide detailed predictions of atomic and molecular emission and absorption that track the formation of a molecular cloud from a purely atomic medium, with a view toward helping to verify cloud formation by shock waves. However, our predictions suggest that the detection of the pre-CO stages will be challenging. Finally, we provide an analytic solution for time-dependent $\\h2$ formation which may be of use in numerical hydrodynamic calculations. ", "introduction": "The formation of stars and planetary systems is one of the fundamental problems in astrophysics. Much of the work over the past decades has examined the formation of low mass stellar systems because these objects sometimes form in isolation and are therefore easier to study individually. One of the earliest stages of stellar birth that has been the focus of numerous investigations is the creation of a centrally concentrated molecular core from a portion of a low-density parent giant molecular cloud (GMC). Theoretical models account for the condensation as occurring possibly via the slow diffusion of magnetic flux occurring over long timescales ($\\sim 10$ Myr) (Mouschovias 1999; Lizano \\& Shu 1989) or the dissipation of turbulence on shorter timescales (Stone, Ostriker, \\& Gammie 1998; Mac Low \\etal 1998; Myers \\& Lazarian 1998; Nakano 1998). Observations of isolated pre-stellar molecular cores, such as L1544 in the Taurus Molecular Cloud, have provided fertile ground for comparison to these theories (Caselli et al 2002; Ciolek \\& Basu 2000; Williams et al 1999; Tafalla et al 1998). However, it is now recognized that most stars form in groups -- from small aggregates to large clusters -- and it is not clear that all theories developed for isolated star formation are easily applicable to the larger scales and simultaneity required for star cluster formation (Ballesteros-Paredes, Hartmann, \\& V\\'azquez-Semadeni 1999). Expanding to larger scales opens the question as to whether the formation of stars, both isolated and clustered, might perhaps be intimately related to the formation of the GMC itself. In a previous paper (Hartmann, Ballesteros-Paredes, \\& Bergin 2001 $=$ HBB01) we pointed out that the great majority of molecular cloud complexes in the solar neighborhood appear to be forming young stars, and that the ages of the stellar populations in these clouds are typically $\\sim$~2~Myr; stellar associations of ages $\\simgreat$~10 Myr are devoid of molecular gas. The calculations indicating that MHD turbulence damps rapidly (Stone et al. 1998, Mac Low et al. 1998, Padoan \\& Nordlund 1999) also favor the suggestion of short cloud lifetimes, since there is no need for a continuous regeneration of MHD turbulence; additional support is found through arguments related to to cloud crossing times (Elmegreen 2000; HBB01). These results place significant empirical constraints on the mechanism(s) of nearby molecular cloud formation \\citep[see][for a review]{maclow_rvmp, balles_cform, elmegreen_araa}. For example, it is difficult to reconcile for the transient nature of local clouds with models in which complexes are built up by the coalescence of smaller molecular clouds, or theories in which molecular gas is mostly moved around from one place to another (see Elmegreen 1993 and references therein), because these processes are likely to take much more than $\\sim 10-20$ Myr to occur. Moreover, short cloud lifetimes also place severe constraints on the processes leading toward fragmentation of the GMC and the condensation towards star formation. In HBB01 we suggested that chemical transformations of local gas are essential to understanding the observational constraints. We suggested that most clouds are formed by large scale flows in the diffuse atomic medium, and that they appear as molecular clouds only when the column density becomes high enough to shield the molecular gas from the dissociating effects of the interstellar radiation field (ISRF). We further noted that self-gravity is likely to become important for column densities comparable to that needed for shielding (also Franco \\& Cox 1986), which would explain the rapid onset of star formation after molecular cloud formation. Finally, we suggested that dispersal of star-forming gas is accompanied by a reduction in shielding, so that the gas may revert to an atomic state some time before it is completely physically removed from the neighborhood. These suggested chemical transformations lessen, but do not necessarily eliminate, the difficulty of making clouds ``fast enough''. The formation of H$_2$ from atomic gas, which generally must precede the formation of CO, is not instantaneous. In addition, while H$_2$ formation places important constraints on the problem, it must also be examined in the context of a model that incorporates the effects of H$_2$ (and CO) self-shielding from the ultraviolet (UV) radiation field with extinction by dust grains, each with its own associated timescale. Following the molecular evolution is critical to an understanding of cloud formation because of observational bias; molecular clouds are essentially defined not by H$_2$ emission but through CO emission. In addition, possible atomic precursors are difficult to identify against the galactic H I background (Ballesteros-Paredes \\etal 1999). In this paper we explore the formation of molecular gas in plane-parallel shocks, starting with atomic (neutral) (warm) gas. Our model incorporates all the relevant heating and cooling mechanisms appropriate for the interstellar medium (ISM), including chemical processes relevant to the transformation of atomic gas to molecular form. Koyama \\& Inutsuka (2000) considered a similar problem, and did include the important effects of shielding from UV radiation, but only examined maximum column densities of standard H~I clouds ($\\sim 10^{19} - 10^{20}$ \\cc ). Thus, they only achieve molecular hydrogen fractions of a few percent, insufficient to follow cloud formation. We explore a range of shock velocities comparable to the flows expected in the diffuse atomic medium (Ballesteros-Paredes, Hartmann, \\& V\\'azquez-Semadeni 1999, and references therein); the parameter ranges are also appropriate for models in which GMC formation is induced by galactic spiral density waves or energetic supernova. Moreover, by following the primary cooling lines in the atomic shock and the post-shock evolution we can predict fluxes for key transitions of C II, C I, O I, CO and other species spanning a range of initial ram pressures. This should aid in the search for the progenitors of molecular clouds. Finally as a by-product of this work we provide an analytic solution for the formation of molecules behind a shock and discuss the relation between gas temperature and extinction, both of which may be of some use to MHD modeling. In \\S 2 we describe the model with results and a parameter study provided in \\S 3. In \\S 4 we outline the observational possibilities for the detection of forming molecular clouds via both emission and absorption lines. Section 5 summarizes the implications of these results for star formation in the local neighborhood and in \\S 6 we present our conclusions. ", "conclusions": "We have presented a detailed coupled physical and chemical model that examines the shock compression of atomic gas and the slow transition to molecular form. This model includes all of the physical and chemical processes relevant to investigate the formation of molecular clouds via shocks induced by cloud-cloud collisions, spiral density waves, or supernovae within the dynamic interstellar medium. Our principle results are as follows: 1) We find that the molecular cloud formation timescale is not controlled by the formation rate of H$_2$ on grains. Rather the shielding of molecules from the UV radiation is the limiting parameter. For all but the weakest shocks we find that H$_2$ self-shields quite efficiently. However, CO formation requires shielding of the interstellar radiation by dust grains. Thus the cloud formation timescale is effectively set by the time needed to accumulate a column equivalent to A$_V \\sim 1$ mag in extinction. 2) Molecular cloud formation times can be as short as $\\sim 10 - 20$ Myr, as required by our picture of rapid cloud formation from large scale flows, adopting typical velocities in the ISM ($v \\sim 10 \\kms$) for starting densities and ram pressures a few times higher than average interstellar values. 3) Since shielding is required for CO formation the a priori presence of H$_2$ in the low density medium will not appreciably shorten the time required to create molecular clouds. 4) We provide detailed predictions of the atomic and molecular emission and absorption that trace the formation of molecular clouds. A subset of these predictions match current conditions observed in over-pressurized gas within the cold neutral medium of the galaxy by Jenkins \\& Tripp (2001). 5) A by-product of this work is an examination of ways to incorporate the effects of chemistry into detailed MHD simulations of structure formation in the galaxy. This includes an analytic solution for the time-dependent formation of molecular hydrogen and a discussion of the overall temperature structure of dense cooling proto-clouds." }, "0405/astro-ph0405435_arXiv.txt": { "abstract": "We present observations of the CO 2-1 and 3-2 transitions toward the merging galaxies of NGC6090 with the Submillimeter Array (SMA)\\footnote[5]{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.}. The high resolution CO data reveal three gas concentrations. The main component is peaking in the overlap region between the two galaxies, where the near-IR and radio continuum emission are weak. The CO 2-1 emission from the face-on galaxy NGC6090E is somewhat stronger than that from the edge-on galaxy NGC6090W. The CO 3-2 emission peaks in the overlap region, similar to the CO 2-1 emission. More than 50\\% of the CO 3-2 emission arises from the 2$''$ (1.2 kpc) area of the overlap region. There appears to be CO 3-2 emission toward the nuclear region and the north-west arm of NGC6090E, while no CO 3-2 emission is detected toward NGC6090W. Unlike the CO gas, most of the radio continuum emission comes from NGC6090E. The strong CO emission, together with the weak radio continuum emission, suggests that star formation in the overlap region has not proceeded long enough to produce significant numbers of supernovae which would be detectable due to their radio continuum emission. ", "introduction": "Galaxy-galaxy interactions and mergers may trigger starburst and nuclear activity, and produce luminous/ultra-luminous infrared galaxies. The study of the kinematics and the distribution of molecular and atomic gas at high angular resolution can help us to understand how the gas responds to merging and interacting processes. By studying different transitions of CO emission, one can determine the excitation conditions of the molecular gas. Similarly, measurements of millimeter/sub-millimeter continuum emission allow us to study the properties of dust. High-J CO transitions, such as CO 3-2 which traces the warm ($\\sim$30K) and dense ($\\sim$10$^4$cm$^{-3}$) molecular gas, pinpoint sites of on-going star formation better than the low-J transitions. With these objectives in mind, we mapped a gas rich merger NGC6090 with the partially completed SMA in the CO 3-2 and CO 2-1 transitions. NGC6090 is a nearby IR-luminous galaxy-galaxy merger, at a distance of 122 Mpc (using $H_0=75$ km s$^{-1}$ Mpc$^{-1}$, Dinshaw et al. 1999), with L$_{IR}(8-1000 \\mu$$m)\\simeq$$3\\times10^{11} L_\\odot$ (Dinshaw et al. 1999), M$_{H_2}\\simeq$$3.0{\\times}10^{10} M_\\odot$ (Bryant \\& Scoville 1999), and M$_{HI}\\simeq$$1.4{\\times}10^{10}M_\\odot$ (van Driel et al. 2001). It consists of two well separated nuclei seen at optical, near-IR, and radio wavelengths. The projected separation is 5$''$.4 (3.2 kpc) measured at radio wavelengths (Dinshaw et al.1999). There appears to be no evidence at the optical and radio wavelengths for an AGN, but there is clear evidence for starburst activity (Dinshaw et al.1999; Bryant \\& Scoville 1999). CO 1-0 observations (Bryant \\& Scoville 1999) show that the dominant component of molecular gas peaks in the overlap region between the two galaxies, where the near-IR emission and radio continuum are weak (Dinshaw et al. 1999). Single dish CO 2-1 and CO 3-2 observations with the CSO (Glenn \\& Hunter 2001) showed that the CO 2-1 and CO 3-2 are strong enough to be mapped with the SMA. In order to understand the physical properties of the molecular gas, we made high resolution images with the SMA in both the CO 2-1 and CO 3-2 transitions. In Section 2, we describe instrumental parameters of the observations; in Sections 3 and 4, we present the results and compare this merger with other interacting systems; and in Section 5, we present a brief summary. ", "conclusions": "" }, "0405/astro-ph0405573_arXiv.txt": { "abstract": "V405~Aurigae is an intermediate polar showing a double-peaked pulsation in soft X-rays and a single-peaked pulsation in harder X-rays. From \\emph{XMM-Newton\\/} observations we find that the soft band is dominated by blackbody emission from the heated white-dwarf surface. Such emission is at a maximum when either magnetic pole points towards us, explaining the double-peaked pulsation. The symmetry of the pulses requires that the angle between the magnetic and spin axes be high. The single-peaked pulsation in harder X-rays is explained in the usual way, as a result of opacity in the accretion curtains. However, the high dipole inclination means that the accretion curtains are nearly in the plane. Thus the outer regions of the curtains do not cross the line of sight to the accretion footprints, explaining the absence of the deep absorption dip characteristic of many intermediate polars. The sawtooth profile of this pulsation requires that the magnetic axis be offset from the white-dwarf centre. We remark also on the double-peaked optical emission in this star. We suggest that the difference between V405~Aur's spin pulse and those of other intermediate polars is the result of its high dipole inclination. ", "introduction": "\\label{sec:intro} V405 Aurigae (RX\\,J0558.0+5353) was discovered in the \\emph{Rosat} All-Sky Survey and identified as an intermediate polar (a cataclysmic variable with a magnetic white-dwarf primary) by Haberl \\etal (1994). It is notable, firstly, for showing a soft blackbody component in the X-ray spectrum, one of a number of such objects discovered with \\emph{Rosat}. Secondly, its soft-X-ray and optical emission shows a double-peaked modulation at the white-dwarf spin period (e.g.\\ Allan \\etal1996), whereas most of these stars show a single-peaked modulation (see, e.g., Patterson 1994 or Hellier 2001 for reviews of this class). The hard X-ray emission in intermediate polars (IPs) originates below a stand-off accretion shock near the magnetic poles of the white dwarf. The soft blackbody emission is then understood as arising from heated white-dwarf surface around the accretion footprints. This is nearly always seen in the AM~Her class of cataclysmic variable, but it is seen only in some IPs, for which the reason is unclear. The issue of why some IPs show a single-peaked pulsation, whereas others show a double-peaked pulsation, is also unclear. One idea (e.g.\\ Hellier 1996; Allan \\etal1996; Norton \\etal 1999) notes that IPs with shorter spin periods will have smaller magnetospheres in which the accretion discs are disrupted nearer the white dwarf. This could result in shorter, fatter `accretion curtains' of material which might have lower opacity in the vertical direction, thus preferentially beaming X-rays along magnetic field lines. The two magnetic poles would combine to produce a double-peaked pulsation. With longer spin periods, where disc disruption occurs further out, the opposite might hold, with tall, thin accretion curtains preferentially beaming X-rays out of the sides. The two poles would then act in phase, producing a single-peaked pulsation. The \\emph{XMM-Newton\\/} X-ray satellite has a larger collecting area and better spectral resolution than \\emph{Rosat\\/}, allowing us to return to V405~Aur with better X-ray data than previously obtained. We report here on a 30-ks \\emph{XMM-Newton\\/} observation aimed at understanding the pulsation at the 545-s spin period of V405~Aur. ", "conclusions": "\\label{sec:conc} \\emph{XMM-Newton\\/} observations of V405~Aur confirms that the soft X-rays are dominated by a blackbody component, presumably from heated white-dwarf surface surrounding the accretion regions. The heated region covers \\til8\\tim{-4} of the white dwarf, comparable to the upper limit found from eclipse timings of XY~Ari. We propose that the modulation of this component results from foreshortening of the blackbody emitting regions. We view the heated surface most favourably when either pole points towards us, so we see a double-peaked modulation in soft X-rays. The equality of the two maxima requires that the magnetic axis be highly inclined from the spin axis. The hard X-ray emission shows a single-peaked, sawtooth modulation. This does not show the strong energy dependence of absorption, as usual in IPs. We suggest that this is because, with the high dipole inclination, the outer parts of the accretion curtains never cross the line of sight. However, electron scattering and opacity in the highly ionized post-shock column causes the intensity variation with spin phase. The sawtooth shape of the pulsation requires that the magnetic axis be offset from the white-dwarf centre. We also suggest that the high dipole inclination is responsible for the double-peaked optical pulse. The high dipole inclination appears to be the main reason for the differences between V405~Aur's pulsation and typical IP behaviour." }, "0405/astro-ph0405090_arXiv.txt": { "abstract": "{ \\object{EK~Cephei} (HD~206821) is a unique candidate to test predictions based on stellar evolutionary models. It is a double-lined detached eclipsing binary system with accurate absolute dimensions available and a precise determination of the metallicity. Most importantly for our work, its low mass (1.12 $M_{\\sun}$) component appears to be in the pre-main sequence (PMS) phase. We have produced detailed evolutionary models of the binary \\object{EK~Cep} using the CESAM stellar evolution code (Morel \\cite{morel97}). A $\\chi^2$-minimisation was performed to derive the most reliable set of modelling parameters (age, $\\alpha_{\\rm A}$, $\\alpha_{\\rm B}$ and $Y_{\\rm i}$). We have found that an evolutionary age of about 26.8 Myrs fits both components in the same isochrone. The positions of EK~Cep~A and B in the HR diagram are consistent (within the observational uncertainties) with our results. Our revised calibration shows clearly that \\object{EK~Cep~A} is in the beginning of the main sequence, while \\object{EK~Cep~B} is indeed a PMS star. Such a combination allows for a precise age determination of the binary, and provides a strict test of the modelling. In particular we have found that the definition of the time step in calculating the PMS evolution is crucial to reproduce the observations. A discussion of the optimal time step for calculating PMS evolution is presented. The fitting to the radii of both components is a more difficult task; although we managed to do it for EK~Cep~B, EK~Cep~A has a lower radius than our best models. We further studied the effect of the inclusion of a moderate convective overshooting; the calibration of the binary is not significantly altered, but the effect of the inclusion of overshooting can be dramatic in the approach to the main sequence of stars with masses high enough to burn hydrogen through the CNO cycle on the main sequence. ", "introduction": "It is very well known that the interior structure of an evolved star is fixed by the initial mass $M_{\\star}$, the initial chemical composition (initial abundances in hydrogen, helium and metals, $X_{\\rm i}$, $Y_{\\rm i}$ and $Z_{\\rm i}$, respectively) and the age $t_{\\star}$. So, except for the Sun, modelling a single star using stellar evolutionary models on the HR diagram is not a closed problem because the number of parameters to be determined is larger than the observational constraints, the present observed luminosity and effective temperature. This difficulty is increased due to uncertainties on the physics of the models used to describe the stellar interior/external structure. Currently these uncertainties are modelled by free parameters such as the diffusive and mass loss coefficient or the mixing length and overshooting parameters for the convection in the MLT approximation (Mixing Length Theory). Some stellar multiple systems are in excellent position to avoid these difficulties. Good determinations of stellar masses are possible for the components of a given binary. Based on the reasonable assumption of a common origin for both components (which yields the same initial chemical composition and age), the problem of modelling both stars of a binary system is reduced to the determination of three so-called stellar modelling parameters, namely, $Y_{\\rm i}$, $Z_{\\rm i}$, $t_{\\star}$ (since $X{+}Y{+}Z{=}1$ we do not need to determine $X_{\\rm i}$ independently). Moreover, if the stars are young enough (so that diffusion processes have not had time to change the surface composition significantly) and the present surface metallicity of the stars is also known (and is, therefore, equal to the initial metallicity), then $(Z/X)_{\\rm i}$ is known. So, the number of observables (five: the effective temperatures and luminosities for both stars plus the metallicity of the system) is higher than the number of modelling parameters. This allows us to include unknown parameters related to the description of the interior physics. For low mass stars, without mass-loss or strong rotation, the convection is the main source of uncertainties in the description of the stellar interior. So we choose to include the mixing length parameter for each component of a binary system $\\alpha_{\\rm A}$ and $\\alpha_{\\rm B}$. A number of binary systems have been calibrated in this way, e.g. \\object{$\\alpha$~Cen} (Noels \\etal\\ \\cite{noels91}; Morel \\etal\\ \\cite{morel00b}) and \\object{$\\iota$~Peg} (Morel \\etal\\ \\cite{morel00a}); all these stars are, however, main sequence (MS) or post-main sequence stars. A good candidate binary system for calibration must have well-determined luminosities, effective temperatures, metallicities and dynamical masses. Pre-main sequence (PMS) binaries that have all the required characteristics are very rare at present (see Palla \\& Stahler \\cite{palla01}, for some examples of PMS binaries, and Lee \\etal\\ \\cite{lee94}, for lithium abundances for some of them). \\begin{table} \\caption{Properties of \\object{EK~Cep} (HD~206821). All data except chemical composition are from Andersen (\\cite{andersen91}) and Popper (\\cite{popper87}); chemical compositions are from Mart\\'{\\i}n \\& Rebolo (\\cite{martin93}). See also Tomkin (\\cite{tomkin83}), Ebbighausen (\\cite{ebbighausen66}), and Hill \\& Ebbighausen (\\cite{hill84}).}\\label{tab:obs} \\begin{center} \\begin{tabular}{lcc} \\hline \\noalign{\\smallskip} & EK~Cep~A & EK~Cep~B \\cr \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} $M_{\\star}/M_{\\sun}$ & 2.029$\\pm$0.023 & 1.124$\\pm$0.012 \\cr $\\log T_{{\\rm eff},\\star}$ (K) & 3.954$\\pm$0.010 & 3.756$\\pm$0.015 \\cr $\\log \\left( L_{\\star}/L_{\\sun} \\right)$ & 1.17$\\pm$0.04 & 0.21$\\pm$0.06 \\cr $R_{\\star}/R_{\\sun}$ & 1.579$\\pm$0.007 & 1.315$\\pm$0.006 \\cr $\\log g$ (cgs) & 4.349$\\pm$0.010 & 4.251$\\pm$0.006 \\cr Sp. Type & A1.5V & G5Vp \\cr [Fe/H] & ... & +0.07$\\pm$0.05 \\cr \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\end{table} The double-lined eclipsing binary system \\object{EK~Cep} (HD~206821) has these unique characteristics. Accurate absolute dimensions are available (Andersen \\cite{andersen91}; Popper \\cite{popper87}) and a determination of the surface metallicity was made (Mart\\'\\i n \\& Rebolo \\cite{martin93}). Most importantly, the radii of the two stars are much closer than expected given their mass ratio. This indicates that EK~Cep~B is still contracting towards the main-sequence. At the same time, the higher mass of the primary implies that the star must be already in its main sequence evolution. All these make \\object{EK~Cep} a perfect candidate to test theoretical models of pre-ZAMS solar-type stars. A summary of the observational characteristics of this binary is given in Table~\\ref{tab:obs}. Several authors have tried to model this system. A problem frequently found is that EK~Cep~A has a smaller radius (or that EK~Cep~B has a bigger radius) than that predicted by the models, making it difficult to fit both components to the observations in the same isochrone. Claret \\etal\\ (\\cite{claret95}) have computed models that agree reasonably well with the observations; however, the radius of EK~Cep~B is somewhat higher than their models predict (a problem also found in the models used for comparison with the observations in Mart\\'\\i n \\& Rebolo \\cite{martin93}). Y\\i ld\\i z (\\cite{yildiz03}) computed models using a higher metallicity than Claret \\etal\\ and a fast rotating core for EK~Cep~A in order to lower its luminosity and radius; this way, models that fit the observed radii and luminosities of the EK~Cep system can be computed. The metallicity used in Y\\i ld\\i z (\\cite{yildiz03}; $Z\\sim 0.04$) is much higher, however, than that given in Mart\\'\\i n \\& Rebolo (\\cite{martin93}; $Z \\sim 0.02$). The values for the iron abundance given by these authors spread somewhat, making a very precise determination of this abundance difficult; nothing in these values comes close, however, to the value of metallicity used in Y\\i ld\\i z (\\cite{yildiz03}). Our atempt to model this system focuses, instead, on the physical input that might affect the PMS evolution. Our own PMS models, computed using the CESAM stellar evolution code (Morel \\cite{morel97}), suggested that there are effects in the PMS evolution due to the physical inputs (overshooting, time step) that have been ignored or too quickly dismissed in the previous works. ", "conclusions": "EK~Cep is an excellent candidate to test stellar evolutionary models because it has one member in the MS phase and the other in the PMS phase. This difference in the evolutionary phase of the components results in a better calibration of the binary age of 26.8~Myrs. We have shown that an incorrect treatment of the time step for the evolution yields models with a lower luminosity during the PMS phase. If both stars are in the PMS phase, a large time step can be compensated for by a higher initial helium abundance (to increase the luminosity of the models; see Fig.~\\ref{fig:var}). This cannot be done if the stars are in different evolutionary phases because the increase in the initial helium abundance will make both stars brighter, while only the PMS component has a lower luminosity, due to the larger time step. The dependence of the calibration on this key aspect of the modelling might explain the results obtained by Claret \\etal\\ (\\cite{claret95}) and other results cited in Mart\\'{\\i}n \\& Rebolo (\\cite{martin93}), in particular that the radius (and luminosity) of EK~Cep~B are underestimated by evolutionary models. However, we cannot reproduce the radius of EK~Cep~A with our models. The presence of a fast rotating core (as in Y\\i ld\\i z \\cite{yildiz03}) would improve the situation. The use of a correct time step solves only part of the problem, by increasing the radius of EK~Cep~B and reducing the need for a larger radius of EK~Cep~A. We stress that our models were calculated using the metallicities obtained by Mart\\'{\\i}n \\& Rebolo (\\cite{martin93}); if a new determination of the metallicity is made, it would be possible to discern more clearly the origin of this inconsistency. In the light of the new calibration, it becomes clear that EK~Cep is an excellent binary to test different aspects of the physical processes of the evolution. Given the particular configuration of this system, one component (the primary) is insensitive to the mixing length parameter and age, while providing a very precise constraint on the helium abundance for the system. The effect of the overshooting on the evolutionary tracks of EK~Cep~A, although not relevant for this work, is important near the end of the PMS phase. We will explore the effects of the overshooting during the PMS evolution for various types of stars in a subsequent paper (Marques \\etal\\, in preparation). EK~Cep~B, still in the PMS, provides a very accurate indication of the age of the system, while being independent of the overshooting prescription. Such a combination makes EK~Cep a test case of how the detailed observation of young binaries is of great relevance for improving the modelling of the evolution in this phase." }, "0405/astro-ph0405059_arXiv.txt": { "abstract": "It has recently been established that the filaments seen in the Las Campanas Redshift Survey (LCRS) are statistically significant at scales as large as $70$ to $80 \\, h^{-1} {\\rm Mpc}$ in the $\\delta=-3^{\\circ}$ slice, and $50$ to $70 \\, h^{-1} {\\rm Mpc}$ in the five other LCRS slices. The ability to produce such filamentary features is an important test of any model for structure formation. We have tested the \\lcdm model with a featureless, scale invariant primordial power spectrum by quantitatively comparing the filamentarity in simulated LCRS slices with the actual data. The filamentarity in an unbiased \\lcdm model, we find, is less than the LCRS. Introducing a bias $b=1.15$, the model is in rough consistency with the data, though in two of the slices the filamentarity falls below the data at a low level of statistical significance. The filamentarity is very sensitive to the bias parameter and a high value $(b=1.5)$, which enhances filamentarity at small scales and suppresses it at large scales, is ruled out. A bump in the power spectrum at $k \\sim 0.05 \\, h \\, {\\rm Mpc^{-1}}$ is found to have no noticeable effect on the filamentarity. ", "introduction": "Quantifying the clustering pattern observed in the galaxy distribution and explaining its origin has been one of the central themes in modern cosmology (eg. Peebles 1980). Traditionally, correlation functions have been used for this purpose, with the two-point correlation function $\\xi(r)$ and its Fourier transform, the power spectrum $P(k)$, receiving most of the attention. There now exist precise estimates of $\\xi(r)$ (eg. Tucker et al. 1997, LCRS; Hawkins et al. 2003, 2dFGRS; Zehavi et al. 2002 SDSS) and $P(k)$ (eg. Lin et al. 1996, LCRS ; Percival et al. 2001, 2dFGRS; Tegmark et al. 2003a, SDSS ) determined from several extensive redshift surveys. It is found that the large scale clustering of galaxies is well described by a featureless, adiabatic, scale invariant primordial power spectrum in a $\\Lambda {\\rm CDM}$ cosmological model. We shall, henceforth, refer to this combination of the background cosmological model and the power spectrum as the \\lcdm model. The observed CMBR anisotropies (eg. Spergel et al. 2003, WMAP), and the joint analysis of CMBR anisotropies and galaxy clustering data (eg. Tegmark et al., 2003b), which place very precise constraints on cosmological models, are all consistent with the \\lcdm model. This model is now generally accepted as the minimal model which is consistent with most currently available cosmological data. One of the most striking visual features in all the galaxy redshift surveys e.g., CfA (Geller \\& Huchra 1989), LCRS (Shectman et al. 1996), 2dFGRS (Colless et al. 2001, Colless al. 2003) and SDSS (EDR) (Stoughton, et al. 2002, Abazajian et al. 2003). is that the galaxies appear to be distributed along filaments. These filaments are interconnected and form a network known as the ``Cosmic Web''. The region in between the filaments are voids which are largely devoid of galaxies. This description of the galaxy distribution, based on the morphology of coherent structures observed in redshift surveys, addresses an aspect of the clustering pattern which is completely missed out by the two-point statistics like $\\xi(r)$ and $P(k)$. In the currently accepted scenario of structure formation, the primordial density perturbations are a Gaussian random field which is completely described by the power spectrum, the phases of different modes being random. As structure formation progresses, coherent structures like sheets and filaments are formed through the process of gravitational instability. The phases of different Fourier modes are now correlated, and the power spectrum does not fully describe the statistical properties of the large scale structures. The full hierarchy of N-point statistics can, in principle, be used to quantify all properties of the large-scale structures, but these, being hard to determine at large scales, have not been perceived as the optimum tool for this purpose. Direct methods of quantifying the morphology and topology of large scale coherent features (as we discuss later) have been found to be more effective. The ability to produce large-scale coherent structures like the filaments observed in galaxy surveys is an important test of any models of structure formation. Any such test addresses issues beyond the scope of the two-point statistics like the power spectrum. In this paper we ask the question whether the \\lcdm model is consistent with the filaments observed in galaxy redshift surveys. The analysis of filamentary patterns in the galaxy distribution has a long history dating back to papers by Zel'dovich, Einasto and Shandarin (1982), Shandarin and Zel'dovich (1983) and Einasto et al. (1984) . In the last paper the authors analyze the distribution of galaxies in the Local Supercluster where they find filaments whose lengths increase with smoothing and finally get interconnected into an infinite network of superclusters and voids. The percolation analysis and the genus statistics (Zel'dovich et al. 1982; Shandarin \\& Zel'dovich 1983, Gott, Dickinson \\& Mellot 1986) were some of the earliest statistics introduced to quantify the network like topology of the galaxy distribution. A later study (Shandarin \\& Yess 2000) used percolation analysis to show the presence of a network like structure in the distribution of the LCRS galaxies. The large-scale and super large-scale structures in the distribution of the LCRS galaxies have also been studied by Doroshkevich et al. (2001) and Doroshkevich et al. (1996) who find evidence for a network of sheet like structures which surround underdense regions (voids) and are criss-crossed by filaments. The distribution of voids in the LCRS has been studied by M{\\\" u}ller, Arbabi-Bidgoli, Einasto, \\& Tucker (2000) and the topology of the LCRS by Trac, Mitsouras, Hickson, \\& Brandenberger (2002) and Colley (1997). A recent analysis (Einasto et al. 2003) indicates a super cluster-void network in the Sloan Digital Sky Survey also. Colombi, Pogosyan and Souradeep (2001) have studied the topology of excursion sets at the percolation threshold. The minimal spanning tree (Barrow,Bhavsar \\& Sonoda 1985) is another useful way to probe the geometry of large scale structures. The morphology of superclusters in the PSCz has been studied by Basilakos, Plionis \\& Rowan-Robinson (2001) who find filamentarity to be the dominant feature. On comparing their results with the predictions of different cosmological models, they find that a low density $\\Lambda$CDM model is preferred. These conclusions were further confirmed by Kolokotronis, Basilakos \\& Plionis (2002) using the Abell/ACO cluster catalogue. The Minkowski functionals have been suggested as a novel tool to study the morphology of structures in the universe (Mecke et al. 1994; Schmalzing \\& Buchert 1997). Ratios of the Minkowski functionals can be used to define a shape diagnostic 'Shapefinders' which faithfully quantifies the shapes of both simple and topologically complex objects (Sahni, Satyaprakash \\& Shandarin 1998). Bharadwaj et al. (2000) defined the Shapefinder statistics in two dimensions (2D), and used this to demonstrate that the galaxy distribution in the LCRS exhibits a high degree of filamentarity compared to a random Poisson distribution having the same geometry and selection effects as the survey. This analysis provides objective confirmation of the visual impression that the galaxies are distributed along filaments. In a later paper Bharadwaj, Bhavsar and Sheth (2004) used Shapefinders in conjunction with a statistical technique called Shuffle (Bhavsar \\& Ling 1988) to determine the maximum length-scale at which the filaments observed in the LCRS are statistically significant. They found that the largest length-scale at which filaments are statistically significant is between $70$ to $80 \\, h^{-1}$Mpc, for the LCRS $-3^o$ slice. Filamentary features longer than $80 \\, h^{-1}$Mpc, though identified, are not statistically significant. Such features arise from chance alignments of galaxies. Further, for the five other LCRS slices, filaments of lengths $50 \\, h^{-1}$Mpc to $70 \\, h^{-1}$Mpc were found to be statistically significant, but not beyond. Comparing the filamentarity observed in galaxy redshift surveys against the predictions of different models of structure formation provides a unique method for testing these models. Here we present a method for carrying out this analysis in 2D, and as an example we apply it to the LCRS for which the filamentarity has already been extensively studied ( Bharadwaj et al. 2000; Bharadwaj et al. 2004). In this paper we address the question if the \\lcdm model is consistent with the filaments observed in the LCRS. We have used cosmological N-body simulations to generate different realizations of the galaxy distribution one would expect in the LCRS for the $\\Lambda {\\rm CDM}$ model. The actual and simulated LCRS data were analyzed in exactly the same way using Shapefinders to quantify the degree of filamentarity, and the results were compared to test if the predictions of the $\\Lambda{\\rm CDM}$ model are consistent with the LCRS. The LCRS galaxies may be a biased tracer of the underlying dark matter distribution whose evolution is followed by the N-body simulation. We also consider this possibility, and study how varying the bias parameter effects the network of filaments and voids. Various independent lines of investigation seem to indicate that there may be excess power, in comparison to the \\lcdm model, at scales $k \\sim 0.05 \\, h \\, {\\rm Mpc}^{-1}$. The two dimensional power spectrum for the LCRS (Landy et al. 1996) exhibits strong excess power at wavelengths $\\sim 100 \\, h^{-1}$Mpc. The analysis of the distribution of Abell clusters (Einasto et al. 1997a, 1997b) reveals a bump in the power spectrum at $k=0.05 \\, h \\, {\\rm Mpc}^{-1}$. Also, the recent analysis of the SDSS shows a bump in the power spectrum at nearly the same value of $k$ (Tegmark et al. 2003a). Such a bump would be a deviation from the \\lcdm model and would be indicative of something very interesting happening at large scales. We have considered the possibility that there is such a bump in the power spectrum, and we investigated if the high level of filamentarity observed in the LCRS is indicative of excess power at $k=0.05 \\, h \\, {\\rm Mpc}^{-1}$ in the power spectrum. To present a brief outline of our paper, in Section 2. we present the method of our analysis, in Section 3. we present our results and finally in Section 4. we discuss our results and present conclusions. ", "conclusions": "The \\lcdm model has been found to be consistent with very precise observations of large-scale structures in the universe and their imprint on the CMBR. These tests are limited by the fact that most of them are based on two-point statistics which are not sensitive to the large-scale coherent features like the long filamentary patterns seen in galaxy redshift surveys. It has recently been established, for the LCRS, that the filaments are statistically significant to scale as large as $70$ to $80 \\, h^{-1} \\, {\\rm Mpc}$ (Bharadwaj et al 2004). In this paper we have tested whether the \\lcdm model is consistent with the high level of filamentarity observed in the LCRS. A point to note is that the analysis presented here is restricted to two dimensional sections, and structures which appear as filaments could actually be sections of three dimensional planar structures. This issue whether the 2D filaments are actually also filaments in 3D, though not of crucial importance in the present discussion, is an interesting question which we plan to address in later work. The filamentarity in the unbiased \\lcdm model is lower than in the LCRS. Introducing a positive bias increases the filamentarity at small scales and suppresses the filamentarity at large scales. As the bias is increased, the galaxy distribution gets more concentrated in the high density regions, and the very large scale structures are suppressed. The values of the average filamentarity as a function of the filling factor (Figure 2) are sensitive to the bias, and the quantitative analysis of filamentarity holds the possibility of being a sensitive probe of the bias parameter. We have used two values of the bias parameter and find that a large bias $(b=1.5)$ is ruled out. The simulations with a smaller value of bias $(b=1.15)$ give a good fit to most of the LCRS slices, though in two of the slices ($\\delta=-3^{\\circ}, \\, -45^{\\circ}$) the average filamentarity in the simulations falls somewhat below the LCRS for nearly the entire range of filling factor. The statistical significance of this mild discrepancy ($\\chi^2/\\nu \\sim 1.8$) is not straightforward to interpret, and we adopt the point of view that the \\lcdm model with a small bias is consistent with the filamentarity in the LCRS slices. It is interesting to note that the LCRS $\\delta=-3^{\\circ}$ slice, where the filamentarity is somewhat in excess of the biased \\lcdm model, also has statistically significant filaments extending to scales ($70 - 80 \\, h^{-1} \\, {\\rm Mpc}$) larger compared to the other slices ($50 - 70 \\, h^{-1} \\, {\\rm Mpc}$) (Bharadwaj et al 2004). Analysis of the filamentarity in larger redshift surveys and a better understanding of the statistical properties of the errors in the average filamentarity will, in the future, allow this test to impose more stringent constraints on models of structure formation. A point to note is that the values of bias $b=1.15$ and $1.5$ have no special significance, and have been chosen as two convenient values one representing a modest bias and another a high bias. The value $b=1.15$ is consistent with the analysis of the LCRS power spectrum (Lin et al. 1996) who conclude that the value of the bias in the LCRS is in the range $0.7$ to $1.3$. A point to keep in mind is that the volume limited subsamples analyzed here contain only the very bright galaxies, and the bias is known to increase with the intrinsic luminosity of the galaxies (eg. Wild et al. 2004) There have been speculations that there may be a bump in the power spectrum around the wave number $k \\sim 0.05 \\, h \\, {\\rm Mpc^{-1}}$, and this may have an influence on the filamentary pattern seen in galaxy redshift surveys. We have considered a bump in the power spectrum which enhances the power in the range $0.04$ to $0.06 \\, h \\, {\\rm Mpc^{-1}}$, and we find that this does not have a significant influence on the average filamentarity at any value of the filling factor. The bump in the power spectrum, if it exists, seems to be unrelated to the filamentary patterns seen in redshift surveys. It is interesting to compare our results with some of the other tests which probe models of structure formation beyond the two-point statistics. The bispectrum goes one step beyond the power spectrum, and is sensitive to non-Gaussian features. It has been used to test non-Gaussianity in the primordial power spectrum and determine the bias parameter (eg. Verde at al. 2002, Scoccimarro et al. 2001), but it does not tell us very much about individual, coherent features like filaments. The genus statistics quantifies the topology of the galaxy distribution. Studies of the 2D genus curve for the 2dFGRS (Hoyle et. al 2002a) and the SDSS (Hoyle et. al 2002b) are consistent with the predictions of the \\lcdm model. Schmalzing and Diaferio (2000) have calculated the Minkowski functionals of the galaxy distribution in the nearby universe (the Updated Zwicky Catalogue) and they compare this with \\lcdm N-body simulations. They find that galaxy distribution in the simulated distributions is less coherent than what is observed. Sheth et al. (2003) have developed a method called ``Surfgen'' for generating triangulated surfaces from a discrete galaxy distribution, and have used this to calculate the 3D Shapefinders. They have applied this to a variety of simulated data (Shandarin, Sheth and Sahni 2003, Sheth 2003), but results are awaited from real redshift surveys. In conclusion we note the large-scale coherent features seen in galaxy redshift surveys provides a unique testing ground for probing the models of structure formation beyond the two-point statistics. These tests are also sensitive to the bias and hold the possibility of giving accurate estimates of the bias parameter. The analysis reported in this paper shows the \\lcdm model with a modest bias to be roughly consistent with the filaments observed in the LCRS. Future work using large redshift surveys like the 2dFGRS and the SDSS should be able to accurately quantify the network of large scale coherent structures and place stringent limits on models of structure formation, taking precision cosmology into new grounds beyond the realm of two point statistics." }, "0405/astro-ph0405215_arXiv.txt": { "abstract": "Planets form in the circumstellar disks of young stars. We review the basic physical processes by which solid bodies accrete each other and alter each others' random velocities, and we provide order-of-magnitude derivations for the rates of these processes. We discuss and exercise the two-groups approximation, a simple yet powerful technique for solving the evolution equations for protoplanet growth. We describe orderly, runaway, neutral, and oligarchic growth. We also delineate the conditions under which each occurs. We refute a popular misconception by showing that the outer planets formed quickly by accreting small bodies. Then we address the final stages of planet formation. Oligarchy ends when the surface density of the oligarchs becomes comparable to that of the small bodies. Dynamical friction is no longer able to balance viscous stirring and the oligarchs' random velocities increase. In the inner-planet system, oligarchs collide and coalesce. In the outer-planet system, some of the oligarchs are ejected. In both the inner- and outer-planet systems, this stage ends once the number of big bodies has been reduced to the point that their mutual interactions no longer produce large-scale chaos. Subsequently, dynamical friction by the residual small bodies circularizes and flattens their orbits. The final stage of planet formation involves the clean up of the residual small bodies. Clean up has been poorly explored. ", "introduction": "\\footnotetext{Posted with permission, from the {\\it Annual Review of Astronomy and Astrophysics}, Volume 42 \\copyright 2004 by Annual Reviews, www.annualreviews.org} The subject of planet formation is much too large for a short review. Our coverage is selective. For a broader perspective, the reader is encouraged to peruse other reviews (e.g., \\citen{K02}; \\citen{Lis93}; Lissauer et al. 1995; Lissauer 2004; \\citen{S72}; Wuchterl, Guillot \\& Lissauer 2000). Except for a description of gravitational instabilities in cold disks in Section \\ref{subsec:toomre}, we bypass the crucial stage during which dust grains accumulate to form planetesimals. Our story begins after planetesimals have appeared on the scene. Moreover, we focus on processes that occur in gas-free environments. These were likely relevant during the later stages of the growth of Uranus and Neptune. We devote the first half of the review to the basic physical processes responsible for the evolution of the masses and velocity dispersions of bodies in a protoplanetary disk. Rather than derive precise formulae governing the rates of these processes, we motivate approximate expressions that capture the relevant physics and refer the reader to more complete treatments in the literature (\\citen{BP90}; \\citen{DT93}; Greenberg et al. 1991; Greenberg \\& Lissauer 1990, 1992; \\citen{HN90}; Hornung, Pellat \\& Barge 1985\\nocite{HPB85}; \\citen{IN89}; \\citen{Oht99}; Ohtsuki, Stewart \\& Ida 2002\\nocite{Oht02}; Rafikov 2003a,b,d; \\citen{SI00} ; Stewart \\& Wetherill 1988). Most of the basic physical processes are, by now, well understood as the result of extensive analytical and numerical investigations by many workers. The velocity dispersion in the shear dominated regime is a notable exception, and here we contribute something new. The second half of the review is concerned with the growth of planets starting from a disk of planetesimals. Readers not interested in the derivations of equations describing mass and velocity evolution can skip directly to the second half (beginning with Section \\ref{sec:growthofplanets}). Although many of our results are general, we concentrate on the formation of the outer planets, Uranus and Neptune. We make this choice for two reasons: The formation of Neptune and Uranus presents the most severe timescale problem. Yet it might be the simplest because it was probably completed in the absence of a dynamically significant amount of gas. Order-of-magnitude estimates, particle-in-a-box simulations, and direct N-body simulations have been used to address this problem. We briefly review the weaknesses and strengths of each approach. Then we show how the simplest of these, the two-groups approximation, can capture many of the results obtained in more sophisticated treatments. We show that the evolution of the mass spectrum can be either orderly, neutral, or runaway and describe oligarchy. We conclude with a discussion of how the Solar System evolved from its state at the end of oligarchy to its present state. ", "conclusions": "We begin this review (Sections \\ref{sec:thehillsphere}--\\ref{sec:massgrowthrate}) with an overview of the physics of coagulation: how bodies in a circumsolar disk grow by accreting each other, and how they stir and damp each other's velocities through viscous stirring and dynamical friction. We present order-of-magnitude derivations for relevant formulae that capture the underlying physics. For $uv_H$ and with $u$100 keV) the source contribution is not thought to be critical. At lower energies an adequate \\gray source catalogue (e.g. from INTEGRAL) was not yet available, and it was stated in \\citet{strong03} that point-source contamination was likely. Note that the adopted method of analysis has the property that sources which are unaccounted for are mainly `absorbed' into the fitted instrumental background rather than affecting the diffuse signal. Meanwhile, a new analysis of IBIS data \\citep{lebrun04,terrier04} has shown that indeed a large fraction of the total \\gray emission from the inner Galaxy is due to sources, at least up to 100 keV. This work has produced a source catalogue containing 91 sources, which can now be used as input to the SPI model fitting, giving a much more solid basis for the analysis. This exploits the complementarity of the instruments on INTEGRAL for the first time in the context of diffuse emission. ", "conclusions": "We have exploited the complementarity of the instruments on INTEGRAL to improve the SPI analysis of diffuse Galactic \\gray emission from 20 -1000 keV. Inclusion of the sources found by IBIS improves the sensitivity and reliability of the analysis. Diffuse emission is detected at a level rather lower than previously when the IBIS sources are accounted for. However this marks only the beginning of what will be possible by jointly analysing the data from both INTEGRAL instruments with their unique combination of characteristics. \\begin{figure} \\centering \\includegraphics[width=1.05\\linewidth]{fig1.ps} \\caption{INTEGRAL pointings used for this analysis. \\label{fig:single}} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=1.05\\linewidth]{fig2.ps} \\caption{Instrumental background scaling factor determined by the fitting procedure. Note that the pointing number is sequential, and there is a time gap between the 1st and 2nd GCDEs around pointing 1600. \\label{fig:single}} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=1.05\\linewidth]{fig3.ps} \\caption{Diffuse continuum (blue diamonds) and summed source (red squares) spectra, from this SPI analysis. The diffuse emission is the sum of the three fitted model components. \\label{fig:single}} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=1.05\\linewidth]{fig4.ps} \\caption{Diffuse continuum from this SPI analysis (blue diamonds). Comparison with with 1. OSSE spectrum around $l = b = 0$ from \\citet{kinzer99}, with components: exponentially cutoff power law (red), high-energy continuum power law (blue), positronium (green) and total (black). Components are scaled as in \\cite{strong03} and described in the text. 2. RXTE from \\cite{revnivtsev03} and Revnivtsev (private communication) (light blue bars), 3. COMPTEL from \\cite{strong99} (magenta). \\label{fig:single}} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=1.05\\linewidth]{fig5.ps} \\caption{ Summed SPI source spectra from this analysis (red squares) compared with ridge emission measured by COMPTEL from \\cite{strong99} (magenta). \\label{fig:single}} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=1.05\\linewidth]{fig6.ps} \\caption{ Diffuse Ridge emission measured by SPI (this work): dark blue, compared with results from IBIS \\citep{terrier04}: red. \\label{fig:single}} \\end{figure}" }, "0405/astro-ph0405353_arXiv.txt": { "abstract": "An oscillating universe cycles through a series of expansions and contractions. We propose a model in which ``phantom'' energy with $p < - \\rho$ grows rapidly and dominates the late-time expanding phase. The universe's energy density is so large that the effects of quantum gravity are important at both the beginning and the end of each expansion (or contraction). The bounce can be caused by high energy modifications to the Friedmann equation, which make the cosmology nonsingular. The classic black hole overproduction of oscillating universes is resolved due to their destruction by the phantom energy. ", "introduction": "In this Letter, we consider a scenario in which the universe oscillates through a series of expansions and contractions. After it finishes its current expanding phase, the universe reaches a state of maximum expansion which we will call ``turnaround'', and then begins to recollapse. Once it reaches its smallest extent at the ``bounce'', it will once again begin to expand. This scenario is distinguished from other proposed cyclic universe scenarios \\cite{tolman,steinturok} in that cosmological acceleration due to ``phantom'' energy ({\\it i.e.}, dark energy with a supernegative equation of state, $p < - \\rho$) \\cite{phantom} plays a crucial role. In addition, our work differs from recent proposals in that our model takes place in 3 space and 1 time dimension (though the proposed mechanism for the bounce arises from braneworld scenarios). The idea of an oscillating universe was first proposed in the 1930's by Tolman. Over the subsequent decades, two problems stymied the success of oscillating models. First, the formation of large scale structure and of black holes during the expanding phase leads to problems during the contracting phase \\cite{dickepeebles}. The black holes, which cannot disappear due to Hawking area theorems, grow ever larger during subsequent cycles. Eventually, they occupy the entire horizon volume during the contracting phase so that calculations break down. (Only the smallest black holes can evaporate via Hawking radiation.) The second unsolved problem of oscillating models was the lack of a mechanism for the bounce and turnaround. The turnaround at the end of the expanding phase might be explained by invoking a closed universe, but the recent evidence for cosmological acceleration removes that possibility. For the observationally favored density of ``dark energy'', even a closed universe will expand forever. Thus, cyclic cosmologies appeared to conflict with observations. Our scenario resolves these problems. Our resolution to the black hole overproduction problem is provided by a ``phantom'' component to the universe, which destroys all structures towards the end of the universe's expanding stage. Phantom energy, a proposed explanation for the acceleration of the universe, is characterized by a component $Q$ with equation of state \\begin{equation} w_Q = p_Q/\\rho_Q < -1 . \\end{equation} Since the sum of the pressure and energy density is negative, the dominant energy bound of general relativity is violated; yet recent work explores such models nevertheless. Phantom energy can dominate the universe today and drive the current acceleration. Then it becomes ever more dominant as the universe expands. With such an unusual equation of state, the Hawking area theorems fail, and black holes can disappear \\cite{davies}. In ``big rip'' scenarios \\cite{ckw}, the rapidly accelerating expansion due to this growing phantom component tears apart all bound objects including black holes. (We speculate about remnants of these black holes below.) The phantom energy density becomes infinite in finite time \\cite{caldwell,ckw}. The energy density of any field described by equation of state $w_Q$ depends on the scale factor $a$ as \\begin{equation} \\label{eq:grow} \\rho_Q \\sim a^{-3(1+w_Q)} . \\end{equation} Hence, for $w_Q<-1$, $\\rho_Q$ grows as the universe expands. Of course, we expect that an epoch of quantum gravity sets in before the energy density becomes infinite. We therefore arrive at the peculiar notion that quantum gravity governs the behavior of the universe both at the beginning and at the end of the expanding universe (i.e., at the smallest and largest values of the scale factor). Here we consider an example of the role that high energy density physics may play on both ends of the lifetime of an expanding universe: we consider the idea that large energy densities may cause the universe to bounce when it is small, and to turn around when it is large. The idea is economical in that it is the {\\em same physics} which operates at both bounce and turnaround. In this Letter we use modifications to the Friedmann equations to provide a mechanism for the bounce and the turnaround that are responsible for the alternating expansion and contraction of the universe. In particular, we focus on ``braneworld'' scenarios in which our observable universe is a three-dimensional surface situated in extra dimensions. Several scenarios for implementing a bounce have been proposed in the literature \\cite{shtanov,branebounce}. As an example, we focus on the modification to the Randall-Sundrum \\cite{RSI} scenario proposed by Shtanov and Sahni \\cite{shtanov}, which involves a negative brane tension and a timelike extra dimension leading to a modified Friedmann equation. Another example is the quantum bounce in loop quantum gravity \\cite{singh}. Once the energy density of the universe reaches a critical value, cosmological evolution changes direction: if it has been expanding, it turns around and begins to recontract. If it has been contracting, it bounces and begins to expand. We emphasize that the two components we propose here work together: we use a modified Friedmann equation as a mechanism for a bounce and turnaround, and we add a phantom component to the universe to destroy black holes. Due to the phantom component, the same high energy behavior that produces a bounce at the end of the contracting phase also produces a turnaround at the end of the expanding phase. In addition, the bounce and turnaround are both nonsingular, unlike the cyclic scenario proposed by Steinhardt and Turok \\cite{steinturok}, which is complicated by a number of physical singularities related to brane collisions near the bounce \\cite{singularity}. This is currently a very controversial topic. ", "conclusions": "Our proposal contains the novel feature that both bounce and turnaround are produced by the same modification to the Friedmann equation. However, it does so at the price of including more than one speculative element: the modified Friedmann equation requires a braneworld model to achieve, and the cosmology must be dominated by phantom energy. In many cases a phantom component is difficult to implement from a fundamental standpoint without severe pathologies such as an unstable vacuum (see, for example, Ref. \\cite{Cline:2003gs}.) However, Parker and Raval \\cite{parker} have investigated a cosmological model with zero cosmological constant, but containing the vacuum energy of a simple quantized free scalar field of low mass, and found that it has $w<-1$ without any pathologies. Several additional areas also remain to be addressed. First, as the universe is contracting, those modes of the density fluctuations that we usually throw away as decaying (in an expanding universe) are instead growing. Hence dangerous structures may form during the contracting phase. At the end of the contracting phase, there is no phantom energy to wipe out whatever structure is formed. In this sense, the initial conditions for structure formation in this picture are set either during the phantom energy dominated epoch near turnaround or by the quantum generation of fluctuations in the collapsing phase \\cite{allen}. Black hole formation could still kill the model. Second, it is not obvious that it is possible to create a truly cyclic ({\\it i.e.} perfectly periodic) cosmology within the context of the ``Phantom Bounce'' scenario. The reason for this is entropy production. We speculate that it may be possible to create quasi-cyclic evolution by redshifting entropy out of the horizon during the period of accelerating expansion. Even more speculatively, we note that the special case of $w_Q = -7/3$, although disfavored by observation, possesses an intriguing duality between radiation ($\\rho_{rad} \\propto a^{-4}$) and phantom energy ($\\rho_Q \\propto a^4$). In this case, the behaviors of these components exchange identity under a transformation $a \\rightarrow 1/a$ \\cite{duality}, effectively exchanging bounce for turnaround, a symmetry which might be exploited to achieve truly cyclic evolution." }, "0405/gr-qc0405122_arXiv.txt": { "abstract": "\\vspace*{0.2cm} The late-time tail behavior of massive Dirac fields is investigated in the Schwarzschild black-hole geometry and the result is compared with that of the massive scalar fields. It is shown that in the intermediate late times there are three kinds of differences between the massive Dirac and scalar fields, (I) the asymptotic behavior of massive Dirac fields is dominated by a decaying tail without any oscillation, but the massive scalar field by a oscillatory inverse power-law decaying tail, (II) the dumping exponent for the massive Dirac field depends not only on the multiple number of the wave mode but also on the mass of the Dirac field, while that for the massive scalar field depends on the multiple number only, and (III) the decay of the massive Dirac field is slower than that of the massive scalar field. ", "introduction": "The dynamical physical mechanism responsible for the relaxation of perturbation fields outside a black hole and the decay rates of the various perturbations have been extensively studied \\cite{3}-\\cite{Ching} since Wheeler introduced the no-hair theorem in the early 1970s \\cite{1,2}. The massless neutral external perturbations were first studied by Price, and it was found that the late-time behavior for a fixed ~$r$ is dominated by the factor ~$t^{-(2l+3)}$ for each multiple moment ~$l$ \\cite{3}. The massless charged scalar field was studied in Refs. \\cite{4}-\\cite{6} and the conclusion was that a charged hair decay slower than a neutral one, i.e., the charged scalar hair outside a charged black hole is dominated by a ~$t^{-(2l+2)}$ tail. The massless late-time tail for the gravitational, electromagnetic, neutrino and scalar perturbations had also been considered in the case of the Kerr black holes in Refs. \\cite{7}-\\cite{9}. On the other hand, many authors found that the analysis of massive fields is also physically important since massive fields can cause interesting phenomena which are qualitatively different from the massless case. The evolution of a massive scalar field in the Schwarzschild background was analyzed by Starobinskii and Novikov \\cite{Starobinskii}, and they found that, because of the mass term, there are poles in the complex plane closer to the real axis than in the massless case, which leads to inverse power-law behavior with smaller indices than the massless case. Hod and Piran \\cite{13} pointed out that, if the field mass $\\mu$ is small, namely ~$\\mu M\\ll 1 $, the oscillatory inverse power-law behavior \\begin{eqnarray} \\label{1}&& \\Phi\\sim t^{-(l+3/2)}\\sin(\\mu t), \\end{eqnarray} dominates as the intermediate late-time tails in the Reissner-Nordstr\\\"{o}m background. We \\cite{Jing1} recently investigated the late-time tails of the massless and the self-interacting (massive) scalar fields in a stationary axisymmetric Einstein-Maxwell-dilaton-axion black-hole geometry, and found that the dumping exponents is independent of the rotation parameter and the dilaton of the black hole. Although much attention has been paid to the study of the late-time behaviors of the scalar, gravitational, electromagnetic in static and stationary black-hole backgrounds, however, to my best knowledge, at the moment the late-time evolution of the Dirac fields has not been investigated. The aim of this paper is to study the intermediate late-time tail behavior of the massive Dirac fields in the Schwarzschild black-hole background and to see whether or not special properties exist in this case. The plan of the paper is as follows. In Sec.2 the decoupled massive Dirac equations in the Schwarzschild spacetime are presented. In Sec.3 the black-hole Green's function is introduced by using the spectral decomposition method \\cite{17}. In Sec.4 the intermediate late-time evolution of the Dirac massive fields in the Schwarzschild background is investigated. The section V is devoted to a summary and conclusion. In the appendix we study whether the conclusions might change if the tortoise coordinate is defined in the conventional way. ", "conclusions": "The ``tortoise\" coordinate (\\ref{tor}) is a function of the background geometry and the test field, which is different from the usual tortoise one, $r_*=\\int \\frac{d r}{f}$. Might the conclusions of this paper change if we use usual tortoise coordinate instead of the coordinate (\\ref{tor})? In the appendix we will address this question. The Dirac equations are given by \\cite{Page} \\begin{eqnarray} &&\\sqrt{2}\\nabla_{BB'}P^B+i\\mu \\bar{Q}_{B'}=0, \\nonumber \\\\ &&\\sqrt{2}\\nabla_{BB'}Q^B+i\\mu \\bar{P}_{B'}=0, \\end{eqnarray} where $\\nabla_{BB'}$ is covariant differentiation, $A_{BB'}$ is the electromagnetic field potential, $P^B$ and $Q^B$ are the two-component spinors, and $\\mu $ is the particle mass. In the Newman-Penrose formalism \\cite{Newman} the equations become \\begin{eqnarray}\\label{np} &&(D+\\epsilon-\\rho )P^0+ (\\bar{\\delta}+\\pi-\\alpha )P^1=2^{-1/2}i\\mu \\bar{Q}^{1'},\\nonumber \\\\ &&(\\triangle+\\mu -\\gamma )P^1+ (\\delta+\\beta -\\tau )P^0=-2^{-1/2}i\\mu \\bar{Q}^{0'}, \\nonumber\\\\ &&(D+\\bar{\\epsilon}-\\bar{\\rho} )\\bar{Q}^{0'}+ (\\delta+\\bar{\\pi}-\\bar{\\alpha} )\\bar{Q}^{1'}=-2^{-1/2}i\\mu P^{1},\\nonumber \\\\ &&(\\triangle+\\bar{\\mu} -\\bar{\\gamma} )\\bar{Q}^{1'}+ (\\bar{\\delta}+\\bar{\\beta} -\\bar{\\tau} )\\bar{Q}^{0'}=2^{-1/2}i\\mu P^{0}, \\end{eqnarray} For the Schwarzschild spacetime the null tetrad can be taken as \\begin{eqnarray} &&l^\\mu=(\\frac{r^2}{\\Delta}, ~0, ~0, ~0 ), \\nonumber \\\\ &&n^\\mu=\\frac{1}{2}(1, ~-\\frac{\\Delta}{r^2}, ~0, ~0)\\nonumber \\\\ &&m^\\mu=\\frac{1}{\\sqrt{2} r}\\left(0, ~0, ~1, \\frac{i}{sin\\theta}\\right), \\end{eqnarray} Then, if we set \\begin{eqnarray} &&P^0=\\frac{1}{r}f_1(r,\\theta)e^{-i(\\omega t-m\\varphi)}, \\nonumber \\\\ &&P^1=f_2(r,\\theta)e^{-i(\\omega t-m\\varphi)}, \\nonumber \\\\ &&\\bar{Q}^{1'}=g_1(r,\\theta)e^{-i(\\omega t-m\\varphi)}, \\nonumber \\\\ &&\\bar{Q}^{0'}=-\\frac{1}{r}g_2(r,\\theta)e^{-i(\\omega t-m\\varphi)}, \\end{eqnarray} where $\\omega$ and $m$ are the energy and angular momentum of the Dirac particles, Eq. (\\ref{np}) can be simplified as \\begin{eqnarray} \\label{DD} &&{\\mathcal{D}}_0 f_1+\\frac{1}{\\sqrt{2}}{\\mathcal{L}}_{1/2} f_2=\\frac{1}{\\sqrt{2}}i\\mu r g_1, \\nonumber \\\\ &&{\\Delta \\mathcal{D}}_{1/2}^{\\dag} f_2-\\sqrt{2}{\\mathcal{L}}_{1/2}^{\\dag} f_1=-\\sqrt{2}i\\mu r g_1, \\nonumber \\\\ &&{\\mathcal{D}}_0 g_2-\\frac{1}{\\sqrt{2}}{\\mathcal{L}}_{1/2}^{\\dag} g_1=\\frac{1}{\\sqrt{2}}i\\mu r f_2, \\nonumber \\\\ && {\\Delta \\mathcal{D}}_{1/2}^{\\dag} g_1+\\sqrt{2}{\\mathcal{L}}_{1/2} g_2=-\\sqrt{2}i\\mu r f_1, \\end{eqnarray} with \\begin{eqnarray} &&{\\mathcal{D}}_n=\\frac{\\partial}{\\partial r}-\\frac{i K} {\\bigtriangleup}+2n\\frac{r-M}{\\bigtriangleup},\\nonumber \\\\ &&{\\mathcal{D}}^{\\dag}_n=\\frac{\\partial}{\\partial r}+\\frac{i K} {\\bigtriangleup}+2n\\frac{r-M}{\\bigtriangleup},\\nonumber \\\\ &&{\\mathcal{L}}_n=\\frac{\\partial}{\\partial \\theta}+\\frac{m}{\\sin \\theta } +n\\cot \\theta,\\nonumber \\\\ &&{\\mathcal{L}}^{\\dag}_n=\\frac{\\partial}{\\partial \\theta}-\\frac{m}{\\sin \\theta } +n\\cot \\theta, \\nonumber \\\\ &&K=r^2\\omega.\\label{ld} \\end{eqnarray} It is now apparent that the variables can be separated by the substitutions \\begin{eqnarray} f_1=R_{-1/2}(r)S_{-1/2}(\\theta),\\nonumber \\\\ f_2=R_{+1/2}(r)S_{+1/2}(\\theta),\\nonumber \\\\ g_1=R_{+1/2}(r)S_{-1/2}(\\theta),\\nonumber \\\\ g_2=R_{-1/2}(r)S_{+1/2}(\\theta). \\end{eqnarray} Thus, we have\\cite{Chand} \\begin{eqnarray} \\label{DDD} \\left[\\Delta {\\mathcal{D}}^{\\dag}_{1/2}{\\mathcal{D}}_0-\\frac{i\\mu\\Delta}{\\lambda+i\\mu r}{\\mathcal{D}} _0-(\\lambda^2+\\mu^2 r^2)\\right]R_{-1/2}=0, \\end{eqnarray} and $\\sqrt{\\Delta}R_{+1/2}$ satisfies the complex-conjugate equation. The decoupled equations can then be explicitly expressed as \\begin{eqnarray} \\label{T1} && \\sqrt{\\Delta} \\frac{d }{d r}\\left(\\sqrt{\\Delta} \\frac{dR_{-1/2}}{d r}\\right)-\\frac{i\\mu\\Delta}{\\lambda+i\\mu r}\\frac{d R_{-1/2}}{d r}+P_- R_{-1/2} =0, \\\\ \\label{T2} && \\frac{1}{\\sqrt{\\Delta}} \\frac{d }{d r}\\left(\\Delta^{3/2} \\frac{dR_{+1/2}}{d r}\\right)+\\frac{i\\mu\\Delta}{\\lambda-i\\mu r}\\frac{d R_{+1/2}}{d r}+P_+ R_{+1/2} =0, \\end{eqnarray} with \\begin{eqnarray} && P_-=\\frac{K^2+i (r-M) K}{ \\Delta} -2 i \\omega r -\\frac{\\mu K}{\\lambda+i\\mu r} -\\mu^2r^2-\\lambda^2,\\\\ &&P_+=\\frac{K^2-i (r-M) K}{ \\Delta} +2 i \\omega r+2 s +\\frac{\\mu(i(r-M)-K)}{\\lambda-i\\mu r} -\\mu^2r^2 -\\lambda^2, \\end{eqnarray} where $\\lambda^2=(l-s)(l+s+1)$ is the separation constant. Introducing an usual tortoise coordinate \\begin{eqnarray} r_*=\\int \\frac{r^2}{\\Delta} dr, \\end{eqnarray} and resolving Eqs. (\\ref{T1}) and (\\ref{T2}) in the form \\begin{eqnarray} &&R_{+1/2}=\\frac{\\Delta^{-1/4}}{r}(\\lambda^2+\\mu^2 r^2)^{1/4}e^{-i\\vartheta/2} \\Psi_+,\\nonumber \\\\ &&R_{-1/2}=\\frac{\\Delta^{1/4}}{r}(\\lambda^2+\\mu^2 r^2)^{1/4}e^{i\\vartheta/2}\\Psi_-, \\end{eqnarray} with \\begin{eqnarray} \\vartheta=\\arctan(\\mu r/\\lambda), \\end{eqnarray} we obtain two wave-equations \\begin{eqnarray} && \\frac{d^2 \\Psi_+}{d r_*^2}+\\left\\{\\frac{d H_+}{d r_*}-H_+^2+\\frac{\\Delta}{r^4}P_+\\right\\}\\Psi_+ =0,\\label{LV1}\\\\ && \\frac{d^2 \\Psi_-}{d r_*^2}+\\left\\{\\frac{d H_-}{d r_*}-H_-^2+\\frac{\\Delta}{r^4}P_-\\right\\}\\Psi_- =0, \\label{LV2} \\end{eqnarray} where \\begin{eqnarray} &&H_-=\\frac{1}{4 r^2}\\frac{d\\Delta}{dr}-\\frac{\\Delta}{r^3}+ \\frac{i\\mu}{2(\\lambda+i\\mu r)}\\frac{\\Delta}{r^2}, \\\\ &&H_+=-\\left[\\frac{1}{4r^2}\\frac{d\\Delta}{dr}+\\frac{\\Delta}{r^3} +\\frac{i\\mu}{2(\\lambda-i\\mu r)}\\frac{\\Delta}{r^2}\\right]. \\end{eqnarray} Near the event horizon the asymptotic solutions are \\begin{eqnarray} R_{\\pm 1/2}\\simeq e^{i\\omega r_*}e^{\\mp \\frac{i}{2}tan^{-1}\\left(\\frac{\\mu r}{\\lambda}\\right)},~~~~ or~~~~R_{\\pm 1/2}\\simeq \\Delta^{-s} e^{-i\\omega r_*}e^{\\pm \\frac{i}{2}tan^{-1}\\left(\\frac{\\mu r}{\\lambda}\\right)}. \\end{eqnarray} One may expand the wave-equations (\\ref{LV1}) and (\\ref{LV2}) as a power series in $M/r$ (neglecting terms of order $O\\left(\\left(M/r\\right)^2\\right)$) as follows \\begin{eqnarray} \\label{VV1}&&\\left[\\frac{d^{2}}{dr^{2}}+\\omega^{2}-\\mu^2+ \\frac{4M\\omega^{2}-2M\\mu^2}{r}-\\frac{\\lambda^2+\\frac{\\lambda}{\\mu}\\omega }{r^{2}} \\right] \\xi_\\pm=0, \\end{eqnarray} where $\\xi_{\\pm}=\\left(1-2 M/r\\right)^{1/2}\\Psi_{\\pm}$. It is of interesting to note that Eq. (\\ref{VV1}) becomes Eq. (\\ref{m19}) if we replace $\\lambda$ by $-k$, which shows that the conclusions of this paper do not change if we use usual tortoise coordinate instead of the coordinate (\\ref{tor})." }, "0405/astro-ph0405165_arXiv.txt": { "abstract": "{ {We propose that the mushroom-shaped structure of the Galactic worm GW 123.4--1.5 is created by a cloud collision with the Galactic gas disk.} A hydrodynamic simulation shows that a mushroom-shaped structure is created after the cloud crosses the Galactic midplane. The lifetime of the mushroom-shaped structure is of order the dynamical time scale of the disk, $\\sim 10^7$ years. We find that the velocities across the cap of the mushroom-shaped structure in the simulation are consistent with the observed values. The simulation also predicts a structure on the opposite side of the Galactic plane which is created by the Kelvin-Helmholtz instability after the cloud passes through the disk. ", "introduction": "The Canadian Galactic Plane Survey (Taylor et al. 2003) revealed that the Galactic worm candidate GW 123.4--1.5 was an unusual mushroom-shaped cloud (English et al. 2000). It is hundreds of parsecs in size and unrelated to conventional shell or chimney structure. The mass of the cap is about 4 times greater than that of the stem, and the total mass is estimated to be about $1.55 \\times 10^5 \\Msun$, although such an estimate depends strongly on the distance in which there is some ambiguity. A position-velocity map in the upper portion of the cap shows no velocity gradient. However, the lower portion of the two lobes of the cap that extend back toward the Galactic plane are blueshifted with respect to the central cap region by 5 km/s. One possible origin for the mushroom-shaped cloud is the rise of buoyant gas. This model was studied by English et al. (2000), in the form of a supernova event, and Avillez \\& Mac Low (2001), in the form of bubbles rising from hot gas reservoirs. They showed that a mushroom-shaped cloud was created by the sweeping effect of the buoyant bubble when the bubble rises from the Galactic plane. The gas swept up by the bubble cooled down and increased in density. In this paper, we propose another possible origin of the mushroom-shaped structure: a cloud collision with the Galactic disk. This scenario was first studied by Tenorio-Tagle et al. (1986, 1987). However, no one has yet performed a simulation demonstrating that a cloud impact would generate the mushroom-shaped structure like GW 123.4--15. ", "conclusions": "Based on our simulations, we conclude that a cloud collision with the Galactic disk is one of the possible models of the mushroom-shaped structure, GW 123.4--1.5. The density and column density clearly show a cap and stem structure which resembles the observed structure for both cases of a high-velocity cloud ($v_{zc0}=10 c_{s0} \\simeq 100$ km/s) and an intermediate-velocity cloud ($v_{zc0}=5 c_{s0} \\simeq 50$ km/s). The case of $v_{zc0}=5 c_{s0} \\simeq 50$ km/s shows a better agreement with the size of the mushroom-shaped structure as well as the mass ratio of the cap and stem. It may mean that GW 123.4--1.5 is created by the collision of an intermediate-velocity cloud into the Galactic plane, though the stem in our simulation is a little thicker than that of the observation. However, mushroom-shaped structures may not be commonly observed unless the densities of the impact cloud are comparable to the density of the Galactic plane and the frequency of the impacts is high. In order to get a mushroom-shaped structure, we found that the cloud density should be at least the same order as the disk density. Given a disk number density $\\sim 1$ cm$^{-3}$, the column density of the impact cloud is $> 10^{20}$ cm$^{-2}$. Observations show that such large column densities are not common both for high-velocity clouds (Wakker \\& von Woerden 1997) and intermediate-velocity clouds (Benjamin \\& Danly 1997). Nevertheless, there are some exceptions, such as the Ursa Major intermediate-velocity cloud (Snowden et al. 1994; Benjamin et al. 1996) which has velocity $\\sim 50$ km/s, column density $\\sim 2.0 \\times 10^{20}$ cm$^{-2}$, and estimated size ($\\sim 15 \\times 50$ pc) which implies a number density $\\sim 1$ cm$^{-3}$. The numerical simulations also show that the lifetime of the mushroom-shaped structure is only about a dynamical time, $\\sim 10^7$ years. These results may be some of the reasons that only one mushroom-shaped structure has yet been found near the Galactic plane. The cloud collision model in this paper is an alternative to the buoyant model. The velocity structures across the cap in this model are consistent with observations of GW 123.4--1.5, but it is not obvious that the buoyant model has a similar velocity structure. More detailed kinematical analysis of GW 123.4--1.5 would distinguish the two models. In addition to the velocity, the stem near the midplane of the Galaxy is clear in our model, while it is not clear in the buoyant model. Our model shows that the stem directly connects to the dense gas near the midplane of the Galaxy, which seems to be the same as the observation. Moreover, our model predicts small structures on the opposite side of the Galactic plane (in Figure 2 around $z=-1.5H_0$, and in Figure 4 around $z=-1.0H_0$). These structures may not be significant in current observations, but may be an interesting target in the future in order to distinguish the two models. { We have assumed a Galactic atmosphere with two effective temperatures. The resulting dense sheet of gas near the midplane contributes to the formation of the stem of the mushroom-shaped structure. In our model, the stem originates from material lifted up from the dense sheet, while the cap is formed primarily by the breakup of the impact cloud. A dense stem is actually not clear in single-temperature models that we performed with the same initial impact clouds. A distinct transition between cold (weakly turbulent) and warm (highly turbulent) gases is not obvious from observations which measure the integrated gas along the line of sight. However, the mushroom-shaped structure we have obtained in our simulations may suggest that the cold gas clouds which are located near the midplane of the Galaxy are locally well separated from the warm gas. } Future work can improve upon this model by including a self-consistent treatment of turbulent pressure, and the effect of radiative cooling of the thermal gas. Moreover, the magnetic pressure is comparable to the effective turbulent pressure in the ISM, and a collision exactly perpendicular to the Galactic plane would be rare in a real situation. Three-dimensional simulations including the magnetic field, heating and cooling, and turbulent gas would be an ultimate work for the future in order to have more detailed comparison between the numerical simulations and observations." }, "0405/astro-ph0405243_arXiv.txt": { "abstract": "We present \\ion{H}{1} synthesis maps of the edge-on starburst NGC~5433 and its environment, obtained with the VLA in its C and D configurations. The observations and spectral model residuals of the main disc emission in NGC~5433 reveal 3 extraplanar features. We associate 2 of these features with coherent extraplanar extensions across multiple spectral channels in our data, including a complete loop in position-velocity space. Interpreting the latter as an expanding shell we derive a corresponding input energy of $2 \\times 10^{54}$ ergs, comparable to that for the largest supershells found in the Galaxy and those in other edge-on systems. NGC~5433 is in a richer environment than previously thought. We confirm that KUG~1359+326 is a physical companion to NGC~5433 and find two new faint companions, both with Minnesota Automated Plate Scanner identifications, that we label SIS-1 and SIS-2. Including the more distant IC~4357, NGC~5433 is the dominant member of a group of at least 5 galaxies, spanning over 750~kpc in a filamentary structure. A variety of evidence suggests that interactions are occurring in this group. While a number of underlying mechanisms are consistent with the morphology of the high-latitude features in NGC~5433, we argue that environmental effects may play a role in their generation. ", "introduction": "High-latitude extensions from the interstellar medium (ISM) in spiral galaxies are important probes of disc stability and energetics, halo dynamics, and the physical processes that drive galaxy evolution. These features are best detected in edge-on galaxies, in which extraplanar emission can be clearly separated from the underlying disc. Discrete features (e.g. arcs, plumes, and filaments in various ISM tracers, or `supershells' in \\ion{H}{1} as well as broader-scale structures such as thick discs or haloes) have now been identified in many galaxies using this approach. This has led to ideas of galactic circulation involving `fountains', `chimneys', or possibly outflowing winds which may enrich the intergalactic medium (Shapiro and Field 1976, Bregman 1980, Norman and Ikeuchi 1989, Heckman, Armus \\& Miley 1990, Breitschwerdt, V{\\\"o}lk \\& McKenzie 1991, Breitschwerdt, McKenzie \\& V{\\\"o}lk 1993). Disc-halo features are usually attributed to underlying drivers such as supernovae and stellar winds. However, these processes alone may be insufficient to produce the highly energetic \\ion{H}{1} supershells seen in some galaxies, or to account for the absence of detectable star formation regions or remnants in others (e.g. Rhode et al. 1999, Perna \\& Gaensler 2004). Thus, while star formation may be a key ingredient in the disc-halo phenomenon, it may not, alone, be sufficient to drive disc-halo dynamics. Impacting clouds have long been invoked as a possible alternative (Tenorio-Tagle \\& Bodenheimer 1988, Santill\\'{a}n et al. 1999), but galaxies that are apparently isolated also sometimes display large supershells (King \\& Irwin 1997, Lee \\& Irwin 1997). This has led to speculation about other possible contributors, including gamma-ray bursts (Efremov, Elmegreen \\& Hodge 1998, Loeb \\& Perna 1998), inflated bubbles from radio jets (Gopal-Krishna \\& Irwin 2000) and magnetic fields (Kamaya et al. 1996). Environmental effects have also been invoked to explain a variety of halo phenomena, such as differing extents of radio (Dahlem, Lisenfeld \\& Golla 1995) and X-ray haloes (Wang, Chaves \\& Irwin 2003) among various edge-on systems. This paper represents an attempt to further characterize the importance of environment in driving disc-halo dynamics. A first problem is the lack of information about the environments of galaxies showing disc-halo outflows. We have therefore obtained maps of the large-scale \\ion{H}{1} in the edge-on galaxy NGC~5433 using the Very Large Array\\footnote{The VLA is a facility of the National Radio Astronomy Observatory (NRAO). The NRAO is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} (VLA) in its C and D configurations. We examine the main disc and extraplanar \\ion{H}{1} in NGC~5433 using the higher resolution C configuration data, and combine the C and D datasets to probe its environment. NGC~5433 ($D$ = 65 Mpc; $H_0 = 70\\,\\,\\rm{km\\,s^{-1}\\,Mpc^{-1}}$) is an infrared bright spiral (Soifer et al. 1989) that was observed in the radio continuum surveys of Irwin, English \\& Sorathia (1999; see their fig.~19) and Irwin, Saikia \\& English (2000; see their fig.~1). The 20 cm images from these studies reveal a thick disc which resolves into discrete disc-halo features at high resolution. Single dish \\ion{H}{1} observations of NGC~5433 have been obtained by Mirabel \\& Sanders (1988) and Schneider et al. (1990) and a low sensitivity VLA integrated \\ion{H}{1} map has been presented by Thomas et al. (2002). An optical DSS image of NGC~5433 and its southwest (SW) environment is shown in Fig.~\\ref{fig1}. NGC~5433 appears slightly curved, with an outer ring which is distorted on the north side. Its RC3 (de Vaucouleurs et al. 1991) classification as an Sdm has been revised to an earlier type by Dale et al. (2000), who include NGC~5433 in an {\\it Infrared Space Observatory} Key Project study of normal star-forming galaxies (see also Malhotra et al. 2001). X-ray (Rephaeli, Gruber \\& Persic 1995) and submillimetre (Thomas et al. 2002) detections have also been reported. The basic properties of NGC~5433, obtained from the NASA/IPAC Extragalactic Database (NED) unless otherwise noted, are given in Table~\\ref{basic}. \\begin{figure*} \\includegraphics{fig1_lr.ps} \\caption{DSS image of NGC 5433 and its SW environment. Labels have been placed immediately below the galaxy. The more distant galaxy IC~4357 is not shown; it is located 44 arcmin SW of NGC~5433, along a line joining NGC~5433 and CGCG~191-037. \\label{fig1} } \\end{figure*} Two galaxies in the vicinity of NGC~5433, located 5 arcmin (CGCG~191-037) and 9 arcmin (KUG~1359+326) to the SW, have been identified as possible companions in the Uppsala General Catalogue of Galaxies (Nilson 1973). More recently, the trio have been listed as WBL~485 by White et al. (1999) in their catalogue of nearby poor clusters of galaxies. On larger scales NGC~5433 has been paired with IC~4357, 44 arcmin to the SW and differing in recessional velocity from NGC~5433 by $40~\\rm{km\\,s^{-1}}$, in the Nearby Optical Galaxies (NOG) group assignment of Giuricin et al. (2000). Several other faint galaxies, identified using the Minnesota Automated Plate Scanner (MAPS) by Cabanela (1999) or in the 2 Micron All-Sky Survey (2MASS), are also in the field shown in Fig.~\\ref{fig1}. Two of these have been labelled SIS-1 (MAPS-NGP O\\_325\\_0013717) and SIS-2 (MAPS-NGP O\\_271\\_0297230) for simplicity. There are no galaxies brighter than 17th (blue) magnitude in regions of the same size in other directions with respect to NGC~5433. Available data (from NED) for the galaxies labelled in Fig.~\\ref{fig1} are given in Table~\\ref{basic}. In this paper we examine the \\ion{H}{1} distribution and kinematics of NGC~5433, and characterize its environment for the first time. Our observing and data reduction techniques are given in \\S\\ref{obsred}, and details of the radio frequency interference (RFI) excision we performed are in Appendix~\\ref{rfi}. In \\S\\ref{results} we present our results: an analysis of the \\ion{H}{1} content and kinematics of NGC~5433 is in \\S\\ref{HI_n5433}, and in \\S\\ref{HI_environment}~and~\\S\\ref{HI_companions} we characterize the environment of NGC~5433 and the \\ion{H}{1} in its companions. In \\S\\ref{discussion} we discuss the prevalence of high-latitude \\ion{H}{1} in edge-on spirals, the characteristics of the NGC~5433 group and possible origins for the extraplanar emission observed. A summary of our findings is in \\S\\ref{conclusions}. \\begin{table*} \\begin{minipage}{140mm} \\caption{Basic Properties of NGC 5433 and Nearby Galaxies \\label{basic}} \\begin{tabular}{@{}lcccccccc@{}} \\hline Galaxy & Morph. & $\\alpha$ (J2000) & $\\delta$ (J2000) & Sep.$^a$ & Mag. & Diameter & $V_{sys}$$^b$ & $i$$^c$ \\\\ & Type & ($^{\\rm h}$ $^{\\rm m}$ $^{\\rm s}$) & ($^\\circ$ $\\arcmin$ $\\arcsec$) & (arcmin) & (blue) & (blue; arcmin$^2$) & ($\\rm{km}\\,\\rm{s}^{-1}$) & ($^\\circ$) \\\\ (1) & (2) & (3) & (4) & (5) & (6) & (7) & (8) & (9)\\\\ \\hline NGC~5433 & SAB(s)c:$^d$ & 14 02 36.05 & 32 30 37.8 & ... & 14.10 & $1.6\\,\\times\\,0.4$ & $4354\\,\\pm\\,7$ & 78 \\\\ CGCG 191-037 & ... & 14 02 20.56 & 32 26 53.5 & 5.0 & 15.6 & $0.8\\,\\times\\,0.2$ & ... & 76\\\\ KUG~1359+326 & ...& 14 01 55.87 & 32 27 29.6 & 9.0 & 15.7 & $0.4\\,\\times\\,0.2$ & ... & 60\\\\ SIS-1$^e$ & ...& 14 01 49.62 & 32 20 34.5 & 14.0 & 17.45 & $0.37\\,\\times\\,0.17$ & ... & 63\\\\ SIS-2$^e$ & ...& 14 02 38.97 & 32 27 50.3 & 2.9 & 18.00 & $0.21\\,\\times\\,0.16$ & ... & 40\\\\ IC~4357 & S? & 14 00 43.69 & 31 53 39.0 & 44.0 & 14.75 & $1.1\\,\\times\\,0.6$ & $4394\\,\\pm\\,36$ & 60\\\\ \\hline \\end{tabular} $a)$ Projected angular separation from NGC~5433.\\hfill\\break $b)$ Heliocentric, optical definition (and so throughout).\\hfill\\break $c)$ From Col. 7 assuming a thin disc except for NGC~5433, where an intrinsic axial ratio $q=0.13$ is adopted.\\hfill\\break $d)$ Revised from RC3 (de Vaucouleurs et al. 1991) classification of Sdm, by Dale et al. (2000).\\hfill\\break $e)$ Alternates: SIS-1 = MAPS-NGP O\\_325\\_0013717, SIS-2 = MAPS-NGP O\\_271\\_0297230\\hfill\\break \\end{minipage} \\end{table*} \\section[]{Observations and Data Reduction} \\label{obsred} The data were obtained with the VLA in its C and D configurations on 2001 July 21, 23 and 1999 May 4, respectively. NGC~5433 was at the pointing centre of each observation, and as such all of the galaxies listed in Table~\\ref{basic} except IC~4357 fell within the instrument's primary beam. During each run the galaxy was observed in 25--35 minute intervals, separated by 3--5 minute phase calibrator observations. Online Hanning smoothing was applied. The flux calibrator, also used as a bandpass calibrator, was observed at both the start and end of each observing run. A total of 554 minutes were spent on source in the C configuration, and 142 minutes in the D configuration. Windy conditions throughout the D configuration run caused a number of antennas to stow, reducing the amount of data obtained at those times by an average of 10 per cent. The data from each observing run were calibrated separately using the Astronomical Image Processing System ({\\sevensize AIPS}). Standard flux, phase and bandpass calibration routines were applied, and the continuum emission in the resulting cubes was removed by subtracting linear fits to the real and imaginary parts of the visibility data in the line-free frequency channels of each bandpass. Following these steps, strong RFI was detected in the central 1 or 2 channels of each dataset. The culprit was the 1400 MHz band centre (corresponding to the recessional velocity of NGC~5433) adopted for the observations: with this setup, correlated noise from the 7th harmonic of the VLA's 200 MHz local oscillator at L-band contaminated the shortest east-west baselines in the central channels, for which the expected fringe rate is zero (Bagri 1996). This RFI was excised by {\\it a)} removing baselines with east-west projections close to the array centre {\\em in all channels}, and {\\it b)} clipping contaminated data in the visibility domain from {\\em infected central channels only}. With this approach the RFI was excised at the root mean square (RMS) noise level $\\sigma$ of the infected channels. The main results of the RFI removal are that $\\sigma$ for the central channel of each dataset is $\\sim$10 per cent higher than the corresponding datacube mean, and flux estimates for NGC~5433 from the central 2 channels of the C configuration data have an associated {\\em systematic error} on the order of $\\sigma$; all other flux measurements and HI maps are unaffected. Details of the RFI excision procedure are given in Appendix~\\ref{rfi}. After calibration, the C configuration datasets were combined in the visibility domain and imaged using a variety of weighting schemes to emphasize \\ion{H}{1} structures of different spatial scales. The C and D configuration data were also combined and imaged in a similar manner; we will refer to this combination as C+D data. Each data cube was then {\\sevensize CLEAN}ed (Clark 1980) and corrected for the attenuation of the primary beam. Due to the lower sensitivity (relative to the C+D maps) and resolution (relative to the C maps) of the D configuration data, we show no images of this dataset alone. However, the C and D configuration observations were obtained independently and have similar sensitivities (see Table~\\ref{observe}), and thus the D data provide a useful consistency check on the reality of faint features detected in the C data: one expects to find them in both cubes. All maps displayed are uncorrected for the primary beam, but calculations are performed on the corrected cubes. The observing and map parameters for the C, D and C+D data are given in Table~\\ref{observe}. A map of the continuum emission was also made from the line-free channels and used to measure continuum fluxes in the vicinity of our \\ion{H}{1} detections (Table~\\ref{global_params}). Since the resulting images of NGC~5433 do not improve upon those of Irwin et al. (1999), we do not reproduce them here. \\begin{table*} \\begin{minipage}{125mm} \\caption{\\ion{H}{1} Observing and Map Parameters \\label{observe}} \\begin{tabular}{@{}lccc@{}} \\hline Parameter & C config. & D config. & C+D combined \\\\ (1) & (2) & (3) & (4)\\\\ \\hline Observing date & 2001 July 21 \\& 23 & 1999 May 4 & ---\\\\ On-source observing time $\\,$ (min.) & 554 & 142 & 696\\\\ Band centre $\\,$ (km s$^{-1}$) & 4352 & 4352 & 4352 \\\\ Total bandwidth $\\,$(km s$^{-1}$) & 657 & 657 & 657\\\\ Channel width $\\,$ (km s$^{-1}$) & 10.1 & 21.2 & 21.2 \\\\ Natural weighting: & & & \\\\ $\\,\\,\\,\\,$ Synth. beam $\\,$ (arcsec) $\\,$ @ PA $\\,$ (\\degr) & $19\\,\\times \\, 17$ @ 70 & $53\\,\\times\\,50$ @ -24 & $30\\, \\times \\, 29$ @ -54 \\\\ $\\,\\,\\,\\,$ RMS map noise $\\sigma$$^a$ $\\,$ (mJy $\\mathrm{beam}^{-1}$) & 0.31 & 0.36 & 0.21 \\\\ Robust weighting$^b$: & & & \\\\ $\\,\\,\\,\\,$ Synth. beam $\\,$ (arcsec) $\\,$ @ PA $\\,$ (\\degr) & $16\\times \\, 15$ @ 74 & $47\\,\\times\\,45$ @ -24 & $23\\, \\times \\, 21$ @ -58 \\\\ $\\,\\,\\,\\,$ RMS map noise $\\sigma$$^a$ $\\,$ (mJy $\\mathrm{beam}^{-1}$) & 0.32 & 0.38 & 0.21 \\\\ Uniform weighting: & & & \\\\ $\\,\\,\\,\\,$ Synth. beam $\\,$ (arcsec) $\\,$ @ PA $\\,$ (\\degr) & $13\\,\\times \\,13$ @ 68 & $42\\,\\times\\,39$ @ -26 & $16\\, \\times \\,15$ @ -53 \\\\ $\\,\\,\\,\\,$ RMS map noise $\\sigma$$^a$ $\\,$ (mJy $\\mathrm{beam}^{-1}$) & 0.37 & 0.41 & 0.24 \\\\ \\hline \\end{tabular} $a)$ $\\sigma$ at $V=4352\\,\\,\\rm{km\\,s^{-1}}$ is $\\sim$10 per cent higher than this value; see text and Appendix~\\ref{rfi}.\\hfill\\break $b)$ {\\sevensize AIPS} robustness parameter is chosen to produce a synthesized beam intermediate to natural and uniform weighting.\\hfill\\break \\end{minipage} \\end{table*} \\begin{figure} \\includegraphics{fig2.ps} \\caption{ Global profiles for NGC~5433 {\\it (solid lines)}, KUG 1359+326 {\\it (long-dashed lines)}, SIS-1 {\\it (short-dashed lines)}, and SIS-2 {\\it (dotted lines)}. Solid error bars represent 1$\\sigma$ random errors. For NGC~5433, the dotted part of the error bars represents additional systematic errors on those points from RFI removal; see \\S\\ref{obsred} and Appendix~\\ref{rfi} for details. \\label{fig2} } \\end{figure} ", "conclusions": "\\label{discussion} Our observations probe the \\ion{H}{1} morphology and kinematics of NGC~5433 and its environment with unprecedented detail. In light of our results we discuss the prevalence of high-latitude \\ion{H}{1} in edge-on spirals in \\S\\ref{disc1}, characterize the NGC~5433 system as an interacting group of galaxies in \\S\\ref{disc2}, and underscore the possible importance of environmental effects in driving the disc-halo phenomenon in \\S\\ref{disc3}. \\subsection{High-latitude \\ion{H}{1} in NGC~5433 and other spirals} \\label{disc1} We identify 3 extraplanar features in NGC~5433 by examining moments of the \\ion{H}{1} distribution and best-fitting kinematic models. 2 of these (F1 and F3) appear to be associated with coherent features in P-P and P-V space (Figs.~\\ref{fig3a}, \\ref{fig7}); we find a complete P-V loop in the upper part of feature F3 (Fig.~\\ref{fig7}b), which may represent an expanding shell. For a point injection of energy, the implied input energy for F3 is $E \\sim 2 \\times 10^{54}$ ergs (Equation~(\\ref{injection})). This is nearly an order of magnitude larger than the corresponding value $E=4.6 \\pm 2.0 \\times 10^{53}$ ergs for the most energetic new supershell (GSH 292-01+55) found in the Southern Galactic Plane Survey (McClure-Griffiths et al. 2002), but lower than $E$ inferred for supershells found in other spirals, such as feature F1 in NGC~2613 ($E > 1.4 \\times 10^{55}$ ergs, Chaves \\& Irwin 2001) and F2 in NGC~5775 ($E \\sim 2.1 \\times 10^{55}$ ergs, Lee et al. 2001). Given the order-of-magnitude uncertainties on these estimates (Lee \\& Irwin 1997), it is clear that the implied energetics for the feature F3 in NGC~5433 are comparable to those of the largest structures found in the Milky Way as well as features in other spirals. While the sample remains small, discrete \\ion{H}{1} features and/or thick \\ion{H}{1} discs have now been observed in galaxies with a wide range of star formation rates (e.g. higher: NGC~891, Swaters, Sancisi \\& van der Hulst 1997; NGC~2403, Schaap et al. 2000; low: UGC~7321, Matthews \\& Wood 2003) and in a variety of environments (e.g. isolated: NGC~3044, Lee \\& Irwin 1997; NGC~3556, King \\& Irwin 1997; companions: NGC~5433; interacting: NGC~5775; Irwin 1994). Regardless of its origin, coherent extraplanar \\ion{H}{1} emission may thus be common in spiral galaxies, as is the case for high-latitude cosmic-ray haloes (Irwin et al. 1999). The existence of these features implies significant energy transfers from the disc into the halo, by mechanisms that remain poorly understood. We return to plausible formation scenarios for features in NGC~5433 in \\S\\ref{disc3}. \\subsection{An active environment} \\label{disc2} NGC~5433 is in a richer environment than previously thought, having at least 3 physical companions (KUG~1359+326, SIS-1, SIS-2) rather than the 2 inferred through proximity and size alone (KUG~1359+326, CGCG~191-037). We suspect that other faint MAPS galaxies in the field might also be detectable in \\ion{H}{1} with observations of greater sensitivity. Since we do not detect either \\ion{H}{1} or 20~cm continuum emission from CGCG~191-037, its relationship to the detected galaxies remains unclear; \\ion{H}{1} observations with a wider bandpass would resolve this issue. If the more distant NOG companion IC~4357 is included, the group size increases to 5; \\ion{H}{1} observations of IC~4357 of similar sensitivity to those presented here may also reveal physical companions, making the group richer still. We note that the X-ray detection reported for NGC~5433 (Rephaeli et al. 1995) may represent hot intergalactic gas, sometimes found in poor groups (Fukazawa et al. 2002). It is interesting that all confirmed group members differ in velocity from NGC~5433 by no more than 70 km s$^{-1}$ (Table~\\ref{global_params}). SIS-1, for example, at a projected separation of 248 kpc, is separated in velocity from NGC~5433 by only 25~$\\rm{km\\,s^{-1}}$. Thus the projected separation of the galaxies may be close to their true separation. In addition, within 20 arcmin of NGC~5433 there are over 3 times as many 2MASS galaxies (12 vs. 4) and 2 times as many optically detected galaxies (14 vs. 7) in the NE+SW quadrants than in the NW+SE quadrants. We thus speculate that the NE - SW orientation of the galaxies in this study may represent the largest systems in a ``true\" filament at least $770$~kpc in size (the separation between NGC~5433 and IC~4357), with little variation in depth. Thus the group geometry is in qualitative agreement with galaxy structures on larger scales found in surveys (e.g. Cross et al. 2001) and simulations of hierarchical structure formation via gravitational collapse (e.g. Bertschinger 1998). There is evidence that interactions are taking place in the group, either among the detected galaxies or with a hot intergalactic medium. NGC~5433, itself, shows some optical curvature (Fig.~\\ref{fig1}) and has an asymmetric \\ion{H}{1} distribution (Fig.~\\ref{fig2}), with more \\ion{H}{1} near the close companion SIS-2. All of the detected companions show peculiarities in their \\ion{H}{1} morphology. The channel maps of KUG~1359+326 reveal a large extension that could be tidal in origin (Figs.~\\ref{fig8},~\\ref{fig10}a), and the galaxy's velocity field is highly distorted (Fig.~\\ref{fig10}b). Moment maps of SIS-1 (Figs.~\\ref{fig10}c,~\\ref{fig10}d) show a kinematically distinct feature roughly aligned with the galaxy's major axis, and the \\ion{H}{1} peak in SIS-2 is offset from the optical centre (Fig.~\\ref{fig10}e). Indeed, the better we resolve a companion to NGC~5433, the more distorted its \\ion{H}{1} morphology and kinematics appear. This leads us to suspect that SIS-2 may indeed be strongly interacting with nearby NGC~5433, but that evidence of the encounter is concealed by our poor spatial and spectral resolution of this companion. \\subsection{Underlying mechanisms} \\label{disc3} Since NGC~5433 is a starburst (see fig. 2 of Irwin et al. 1999) it seems plausible that the observed high-latitude extensions are generated by an internal mechanism. A number of scenarios are consistent with the data. Symmetries in features about a galaxy midplane have previously been attributed to a single event in the underlying disc (Chaves \\& Irwin 2001). If this is the case for F2 and F3 (Fig.~\\ref{fig4}a), it is puzzling that we do not find extensions corresponding to F2 in the same channels as those harbouring the shell attributed to F3 (Fig.~\\ref{fig7}b). On the other hand, the double peak in F3 itself and the presence of 20 cm continuum loops in the same vicinity (Fig.~\\ref{fig7}a) suggest a common origin for all of these features; their locations relative to the galaxy nucleus give the impression of a broad-scale nuclear outflow on the west side of the galaxy. Alternatively, the apparent correlation between the continuum and \\ion{H}{1} features is also consistent with a galactic fountain generated by energetic events in the disc (e.g. Lee et al. 2001). None the less, one cannot dismiss the role of NGC~5433's active environment in the generation of the detected extraplanar features. For example, observations of 3 systems with low mass companions by van der Hulst \\& Sancisi (2003) all show evidence for gas accretion on to the parent discs. Conversely, tidal interactions may pull gas out of the midplane of large spirals, sometimes forming arcs and plumes with similar morphologies to supershells (Taylor \\& Wang 2003). In the Galaxy, the origins of some high-velocity clouds (HVCs) have been attributed to these processes (e.g. Putman et al. 2003). We speculate that a similar scenario applies to NGC~5433: given the inferred upper limits on the masses of HVC analogues in external systems (e.g. Pisano et al. 2004), they would not be detectable in our data. Tidal interactions between NGC~5433 and its smaller companions may be responsible for the observed extraplanar emission, either via direct impacts, tidally induced spurs, or by fuelling the increased star formation found in the disc. Additional multi-wavelength observations, particularly high-resolution X-ray data as well as high-sensitivity \\ion{H}{1} and H$\\alpha$ maps, are needed to distinguish between these possibilities. Indeed, the environment can be a key component in the formation of {\\it both} externally and internally generated high-latitude features." }, "0405/astro-ph0405305_arXiv.txt": { "abstract": "{We present a spectroscopic study of D461, a giant star belonging to Draco dwarf spheroidal galaxy. From spectral synthesis in LTE we derive a lithium abundance of $\\rm log$ $\\epsilon$(Li)$=3.5\\pm0.4$ and a C/O ratio between 3 and 5. This is the first detection of a lithium rich C-star in a dwarf spheroidal galaxy. Basing on stellar models of appropriate chemical composition, we show that a similar C enrichment is compatible with that expected for a low mass low metallicity thermally pulsing AGB star, undergoing few third dredge up episodes. The position in the $\\rm log$ g-$\\rm log$ T$_{\\rm{eff}}$ diagram of D461 is also compatible with this theoretical scenario. In particular, the low effective temperature, lower than that expected for a low metallicity giant star, is a consequence of the huge increase of the envelope opacity occurring after the carbon dredge up. The Li enrichment may be explained if a deep circulation would take place during the interpulse period, the so called cool bottom process. In spite of the low resolution of our spectra, we derive a lower limit for the carbon isotopic ratio, namely $^{12}$C/$^{13}$C$>40$, and a constraint for the Ba abundance, namely $0.5<$[Ba/Fe]$< 2$. The proposed scenario also fits these further constraints. Then, we estimate that the mass of D461 ranges between 1.2 and 2 $M_\\odot$, which corresponds to an age ranging between 1 and 3 Gyr. We conclude that this star is more massive and younger than the typical stellar population of Draco. ", "introduction": "The origin of the Galactic lithium is still not completely understood. In addition to that produced by the primordial nucleosynthesis, several sources have been invoked to explain the observed abundances in the various components of the Milky Way (see Travaglio et al. 2001, for a recent analysis of the contributions of various possible producers of galactic Li). Among these, only for asymptotic giant branch (AGB) stars there are observational evidences of some Li production. Li is easily destroyed within stars due to the large cross section of its proton capture reaction. As a consequence, giant stars become Li-depleted after the first dredge up. However, it has been early recognized that a Li production may occur in AGB stars via the beryllium convective belt mechanism (Cameron \\& Fowler 1971). Two are the conditions to be fulfilled: i) the temperature at the base of the convective envelope must be of the order of 20-30 $10^6$ K, so that $^7$Be can be produced via $^3$He$+^4$He reaction, and ii) the mixing must be fast enough to remove the fresh Be from the hot bottom layers, before it decays into Li by electron capture. In such a way, most of the Li will be synthesized in the cool external layers of the star, where the proton capture are defused. Several AGB stars with large abundances of Li are actually found in the Galaxy and in the Magellanic Clouds (Catchpole \\& Feast 1976; Abia et al. 1993; Smith et al. 1995). Li enhancement is commonly found in stellar models of massive AGB stars as a consequence of the hot bottom burning (HBB; Renzini \\& Voli 1981; Sackmann \\& Boothroyd 1992; Forestini \\& Charbonnel 1997; Lattanzio \\& Forestini 1999). However, there are clear indications of an efficient Li production in some low mass AGB stars, where the HBB does not occur. For instance, most of the SLiR \\footnote{We define lithium rich (LiR) and super lithium rich (SLiR) those stars showing log $\\epsilon$(Li)$>1.0$ and log $\\epsilon$(Li)$>4.0$, respectively; as usual, log $\\epsilon$(X) = log (X/H)+12, where X/H is the number of atoms of a given element relative to hydrogen.} AGB stars found in the Galaxy are carbon stars (C/O$>1$ in the envelope, see however Garc\\'\\i a-Lario et al. 1999), while, in the Magellanic Clouds, they are mostly O-rich. Moreover, the luminosities of LiR and SLiR stars in the Clouds are systematically brighter ($-6\\leq \\rm{M_{bol}}\\leq -7$) than their galactic counterparts (M$_{\\rm{bol}}\\geq -5.5$). Their low luminosity and the fact that C/O$>1$ indicate that most of the galactic Li rich stars have low mass progenitors. Finally, the relative abundance of galactic LiR stars with low $^{12}$C/$^{13}$C ratios ($< 15$, mostly C-stars of type J; Abia \\& Isern 1997) is substantially larger than that measured in the Magellanic Clouds (Hatzidimitriou et al. 2003). In summary, the comprehension of the mechanisms of Li production in different AGB stars as well as its dependence on the parent stellar population deserve further observational and theoretical investigations. Dwarf galaxies are other stellar systems where AGB stars can be resolved. In the last few years, a growing amount of AGB stars have been discovered in these galaxies: Withelock et al. (1999) in Sagittarius; Azzopardi et al. (1999) and Demers et al. (2002) in Fornax; Shetrone et al. (2001a) and Margon et al. (2002) in Draco (see also Groenewegen 1999 and references therein). Dwarf galaxies span a wide range in metallicity, so that they could provide new hints about the dependence of the AGB nucleosynthesis (Li production included) on the chemical composition of the parent stellar population. As part of a large survey of AGB stars in the nearby galaxies, we report here, the spectroscopic study of the carbon star D461 in the Draco dwarf spheroidal galaxy, discovered by Armandroff et al. (1995). Let us anticipate our main findings: a) D461 is Li-rich and b)it is more massive and younger than the dominant stellar population of Draco. This is the first report of a star of this type in a dwarf spheroidal. In section 2, we present the observations and the abundance analysis. The evolutionary state of D461 is investigated in section 3, by comparing the available photometric and spectroscopic data to appropriate models of AGB stars. Possible evolutionary scenarios are discussed in the conclusive section. ", "conclusions": "it is a thermally pulsing low mass ($\\sim 1.5$ M$_\\odot$) low metallicity ([Fe/H]$\\sim$ $\\rm{-2)}$ AGB star, undergoing the third dredge up. Gravity and effective temperature are also consistent with this hypothesis. This is the first discovery of an intrinsic C(N) star in a metal poor stellar population. The Li enhancement is likely due to a moderate cool bottom process. This scenario implies that the mass of D461 cannot be smaller than 1.3 M$_\\odot$, because a less massive star would attain the thermally pulsing AGB phase with a too small envelope for the occurrence of the third dredge up. On the other hand, the observed visual magnitude (nearly that of the RGB tip) excludes the possibility that D461 could be a massive AGB star. We estimate that 2 M$_\\odot$ is a conservative upper limit for the mass of D461. Such a stringent limitation of the stellar mass implies a striking constraint to the age of D461: it should range between 1 and 3 Gyr (see e.g. Table 1 in Dom\\'\\i nguez et al. 1999). We remind, again, that the bulk of the stellar population of the Draco dwarf spheroidal is dominated by very old stars ($\\sim 10$ Gyr) although the star formation extended up to 2 Gyr ago. This latter corresponds to the age of a 1.5 M$_\\odot$ AGB star (for the typical metallicities of Draco). Note that the anomalous Cepheids, which are currently detected in Draco, should have masses very close to this value (Bono et al. 1997). Let us finally mention an alternative scenario. It could be possible that D461 was, in origin, a low mass star (M$<$1 M$_\\odot$) whose mass has been recently accreted, as a consequence of a mass exchange or coalescence, from a companion in a close binary system. Up to about 200 Myr ago D461 was an old main sequence star becoming a Blue Straggler and, more recently, a pulsating horizontal branch star with the characteristics of the anomalous Cepheids. Radial velocity monitoring of D461, by Olszewski et al. (1996), found no evidence of the existence of a compact companion, but Shetrone et al. (2001a) detected a possible photometric variability (at 4 $\\sigma$ level). Such a variability is, however, quite common among AGB stars belonging to intermediate-age populations. Additionally, note that the binary hypothesis is disfavoured because of the high Li abundance (see Abia et al. 1993). In summary, the presently available data cannot allow a definitive discrimination between the old and the young hypothesis for D461, even if, from a statistical point of view, the latter is the more significative." }, "0405/astro-ph0405419_arXiv.txt": { "abstract": "{ It is well known that the Solar System is presently moving through a partially ionized local interstellar medium. This gives rise to a counter-flow situation requiring a consistent description of behaviour of the two fluids -- ions and neutral atoms -- which are dynamically coupled by mutual charge exchange processes. Solutions to this problem have been offered in the literature, all relying on the assumption that the proton fluid, even under evidently nonequilibrium conditions, can be expected to stay in a highly-relaxated distribution function given by mono-Maxwellians shifted by the local proton bulk velocity. Here we check the validity of this assumption, calculating on the basis of a Boltzmann-kinetic approach the actually occurring deviations. As we show, especially for low degrees of ionization, $\\xi \\le 0.3$, both the H-atoms and protons involved do generate in the heliospheric interface clearly pronounced deviations from shifted Maxwellians with asymmetrically shaped distribution functions giving rise to non-convective transport processes and heat conduction flows. Also in the inner heliosheath region and in the heliotail deviations of the proton distribution from the hydrodynamic one must be expected. This sheds new light on the correctness of current calculations of H-atom distribution functions prevailing in the inner heliosphere and also of the Lyman-$\\alpha $ absorption features in stellar spectra due to the presence of the hydrogen wall atoms. Deviations from LTE-functions would be even more pronounced in magnetic interfaces, which via CGL-effects cause temperature anisotropies to arise. } ", "introduction": "The problem of the heliospheric interface, where H-atoms and protons are effectively coupled by charge exchange interactions, has often been faced in the literature. In general, it was recognized very early that the passage of neutral interstellar atoms (O, H) through the plasma interface ahead of the solar wind termination shock needs a kinetic treatment, since the relevant charge exchange mean free paths between H-atoms and protons are comparable to or even larger than the typical structure scales of the interface plasma flow (i.e., Knudsen numbers are equal to or smaller than 1; see Ripken \\& Fahr 1983\\nocite{ripken_fahr:83a}; Fahr \\& Ripken 1984\\nocite{fahr_ripken:84}; Fahr 1991\\nocite{fahr:91}; Osterbart \\& Fahr 1992\\nocite{osterbart_fahr:92}; Baranov \\& Malama 1993 \\nocite{baranov_malama:93}; McNutt et al., 1998\\nocite{mcnutt_etal:98}, Bzowski et al. 2000\\nocite{bzowski_etal:00}; Izmodenov et al. 2001) \\nocite{izmodenov_etal:01a}. Nevertheless, due to mathematical complications associated with such kinetic treatments of the problem, many heliospheric models have appeared in the literature which use hydrodynamic treatments of the two fluids, H-atoms and protons, coupled by charge exchange reactions (for recent reviews, see Zank 1999;\\nocite{zank:99} Fahr 2003\\nocite{fahr:03c}b). The hydrodynamics applied in all these approaches is restricted to the description of the space-time behavior of the lowest hydrodynamic moments like density, bulk velocity and pressure and the local distribution function is taken to be solely a function of these moments in the form of shifted Maxwellians. As suggested in Fig.\\ref{figDemo}, the permanent supply of newly charge-exchanged particles into the local distribution functions will maintain the resulting distribution (shown by dashed lines) far from a three-moment HD distribution. This has been clearly recognized by Baranov and Malama (1993, 1995)\\nocite{baranov_malama:93}\\nocite{baranov_malama:95}, who for this reasons perfected a kinetic treatment for the neutral H-atoms, even in view of the mathematical complications. The semi-kinetic approach offered by them, though treating the H-atom kinetically, is still based on the assumption that the protons can be described as a hydrodynamic fluid. But this is not true under special conditions, as we will demonstrate here. The protons generate deviations from a 3-moment HD-distribution. Therefore to represent the two fluids interacting by charge exchange reactions one would need a kinetic description both for the H-atoms and the protons. This highly complicated model will not be offered in this paper, but we present calculations which clearly make visible the resulting deviations of H-atom and proton distributions from 3-moment HD distributions. One other point, which was already emphasized by Fahr (2003a,b),\\nocite{fahr:03a}\\nocite{fahr:03b} concerns the fact that hydrodynamic two-fluid descriptions of the interface flows use charge-exchange coupling terms which are only justified in cases of supersonic bulk velocity differences. Under realistic interface conditions the resulting bulk velocity differences, however, appear to have subsonic magnitudes and thus require revised formulations of the charge exchange coupling terms. In the following we thus subject hydrodynamic approaches to a kinetic control. ", "conclusions": "In purely hydrodynamic multi-fluid interaction codes describing the interaction of the partially ionized interstellar medium with the solar wind plasma, both the proton fluid and the H-atom fluid are generally represented by the moments of distribution functions which are taken to be bulk-flow shifted Maxwellians. We have shown for the region close to the stagnation line, outside of the heliopause but inside of the outer bowshock, that this assumption is substantially violated for the H-atom flow, while it is only mildly violated for the proton flow. This result, however, is due to the values of interstellar parameters adopted in our study, where we have taken the LISM H-atom density = LISM proton density = 0.1~cm$^{-3}$. In that case the p-p Coulomb relaxation processes are quite effective in keeping the resulting proton distribution function close to a shifted Maxwellian. For different LISM parameter combinations, e.g., for lower LISM proton densities, or at different places of the heliospheric interface, the assumption of having the proton distribution function close to a shifted Maxwellian may be violated. This is because the Coulomb relaxation period $\\tau_p \\simeq T_p^{3/2}/n_p$ which means that for lower proton densities and higher proton temperatures this relaxation period may easily increase to values $\\tau_p >> \\tau_{ex}$ and thus induce substantial deviations of the resulting proton distribution function from a shifted Maxwellian. For instance, inside the heliopause, where very low proton densities and high proton temperatures prevail, substantial deviations of the proton distribution from shifted Maxwellians should be expected unless alternative relaxation processes operate, as discussed in Section 4." }, "0405/astro-ph0405133_arXiv.txt": { "abstract": "We present the results of a comprehensive search for stellar variability in the globular cluster 47 Tucanae. Using the Mount Stromlo 40-inch (1m) telescope at Siding Spring Observatory and a combined V+R filter, we have detected 100 variable stars across a 52$\\times$52$'$ field centered on the cluster. The main aim of this project is to search for transiting 'Hot Jupiter' planets in this cluster, the results of which shall be discussed in a separate paper. Here we present the V+R lightcurves and preliminary investigations of the detected variable stars, which comprise 28 Eclipsing Binaries (21 contact binaries and 7 detached systems), 45 RR Lyrae stars (41 of which belong to the Small Magellanic Cloud and four seemingly to the Galactic Halo), and 20 K-giant Long Period Variables (LPVs). We also detected four $\\delta$ Scuti stars, one TypeI Cepheid, and one TypeII Cepheid. One variable appears to be a possible dust-enshrouded SMC star with a short period pulsation. Of these 100 variables, 69 are new discoveries. Our eclipsing binary sample indicates a clear radial segregation in period, and includes two binaries that are seemingly orbited by low-luminosity stellar companions. One RR Lyrae star shows a Blahzko effect with remarkable regularity. Those variables previously known are cross-identified with Kaluzny and coworkers. In agreement with previous studies, this work shows that the relative frequency of detectable variable stars (particularly contact binaries) in 47 Tuc is very low compared to other studied regions. A distance modulus of 18.93$\\pm$0.24 for the Small Magellanic Cloud and 13.14$\\pm$0.25 for 47 Tucanae has been estimated from our sample, and are in agreement with values previously published. The total database presented here contains V and I photometry for 43,067 47 Tuc field stars (13.0$\\leqslant$V$\\leqslant$21.0), along with 33-night V+R lightcurves and astrometry for 109,866 stars (14.5$\\leqslant$V$\\leqslant$22.5). ", "introduction": "Variable stars in globular clusters (particularly binaries) play an important role in understanding cluster dynamical evolution. Despite this, such clusters have seldom been the targets of detailed study for the presence of variables, mainly due to the difficulty in obtaining accurate photometry from the ground of faint stars in very crowded fields. Due to technological advances in recent years, however, the number of clusters studied and the list of variable stars discovered has increased dramatically. As a general overview of the field, those major clusters investigated recently include Omega Centauri \\citep{Kal96,Kal97,Hag03}, M5 \\citep{YR96}, M71 \\citep{YM94}, M4 \\citep{KTK97}, NGC6397 \\citep{K97}, M22 \\citep{Piet03}, NGC6946 \\citep{Guld03}, M69 \\citep{Gregorsok03}, M15 \\citep{Zhel03}, M13 \\citep{Kopacki03} and M3 \\citep{Strader02}. \\citet{YC96} discovered six spectroscopic binaries in NGC5053. \\citet{Clement01} and \\citet{H96} review recent and more historical studies into globular variables. A few previous surveys have searched for variable stars in 47 Tuc. \\citet{Hogg73} discovered two variables. \\citet{Edmonds96a,Edmonds96b} used a HST dataset to detect 75 variable stars, including Eclipsing Binaries and variability among K-giants. \\citet{Alb01} uncovered 107 variable stars, the largest number to date, and derived an overall binary frequency of 14$\\%\\pm$4$\\%$ using the same HST dataset as \\citet{Gil2000} to search for Hot Jupiter planets in the cluster. \\citet{Kal98} performed a wider field survey on the cluster, uncovering 42 variables. We present a new extensive variable star catalogue, a natural byproduct of a photometric project to search for transiting 'Hot Jupiter' planets in 47 Tuc. The results of the planetary search will be the subject of a separate paper. The total catalogue comprises 28 Eclipsing Binary systems, 20 long period variables and 41 Small Magellanic Cloud (SMC) RR Lyraes, four Halo RR Lyraes, two Cepheids and four apparent $\\delta$ Scuti stars, and one anomalous short period Small Magellanic Cloud (SMC) star. Discrimination between cluster/SMC memberships is achieved using the location of the variable on the cluster Colour Magnitude Diagram. Our large number of new discoveries is due primarily to the very large field of view (52$\\times$52$'$) of our survey and also its photometric depth. In this paper, we present the phase-wrapped V+R lightcurves, preliminary investigations into the detected variable stars, and a description of our photometric, astrometric and lightcurve database. Our survey covers a larger area than any previous search, and extends to deeper photometry than presented by \\citet{Kal98}. We recover 31 of Kaluzny's stars and discover a further 69 variables. The unrecovered variables either lie between CCDs, or within regions of no data caused by telescope offsets. The cluster core cannot be easily imaged by ground-based telescopes due to the extreme crowding. On our 300s exposures (V+R), the inner 6$'$ of 47 Tuc is saturated. The cluster field is located at a high galactic latitude (l=305.9 deg, b=$-$44.9 deg), providing low foreground contamination by the Milky Way and low reddening. The field is significantly contaminated by background stars from the Small Magellanic Cloud. Our field of view extends to $\\sim$60$\\%$ of the 47 Tuc tidal radius. Study of the SMC RRLyrae stars, as standard candles, presents an opportunity to investigate the distance to the SMC from a location some seven degrees NW of the centre. Contact eclipsing binaries (EcB) are very useful as distance indicators. Observing any double line spectroscopic binaries offers an opportunity to measure directly the primary stellar parameters such as mass, luminosity, radius and hence distance. \\citet{Ruc93} presents a method of distance determination for contact binaries by including the period, unreddened V-I colour and system metallicity. We present the results of an application of Rucinski's calibration for our binary sample with periods less than 1 day. Section 2 describes the observations and data reduction, along with a description of the method used to obtain the lightcurves. Here we also discuss the astrometry and variable detection methodology, along with notes on cluster membership of the variables. Our survey completeness and the quality of the photometry is described. Section 3 deals with the cluster Colour Magnitude Diagram (CMD), the corresponding photometric database, and describes the photometric calibrations. Section 4 contains a discussion and preliminary analysis of the different types of variables in the catalogue and presents their lightcurves. We summarise and present conclusions in Section 5. Finder charts of these stars are provided on a webpage link to allow easy identification on 47 Tuc wide field images. ", "conclusions": "We have presented data for 100 variable stars detected across a wide (52$\\times$52$'$) field centered on the globular cluster 47 Tucanae. Of these 100, 69 are new discoveries. The sample consists of 41 apparent Small Magellanic Cloud RR Lyrae stars, four Halo RR Lyrae stars, 28 eclipsing binaries, 20 Long Period Variables, four $\\delta$ Scuti Stars and two Cepheids. We also detected one anomalous short period red giant, perhaps surrounded by a dusty region. Four of our RR Lyrae sample display Blahzko Effect variations, one with remarkable regularity. This catalogue more than doubles the number of known variables in the 47 Tuc field. Of the EcB sample, four variables are perhaps orbited by faint companions, most likely M-Dwarf stars. Such stars are important in determining the long-term stability and survivability of low-mass objects in close orbits inside globular clusters. Future spectroscopic observations of these candidates are planned. As well as presenting this new variable catalogue, this paper presents a new complete database of V and I photometry, astrometry and 33-night V+R lightcurves for 109,866 stars across the field. The distance modulus of both 47 Tucanae and the SMC have been determined from our sample. The values are consistent with those already published for these two objects. It is clear from the eclipsing binary results, as well as those previously published in the literature, that the relative frequency of contact binaries in the field of 47 Tuc is very low compared to other studied regions. The reasons for this difference are unclear, although mass-segregation and dynamical effects seem to play an important role. Our sample of EcB shows a distinct period/radial-distance segregation, perhaps indicative of dynamical relaxation. Quite possibly the more massive shorter period contact binaries are located preferentially within the cluster core, an unsampled area in our experiment." }, "0405/hep-th0405054_arXiv.txt": { "abstract": "We explore a dark energy model with a ghost scalar field in the context of the runaway dilaton scenario in low-energy effective string theory. We address the problem of vacuum stability by implementing higher-order derivative terms and show that a cosmologically viable model of ``phantomized'' dark energy can be constructed without violating the stability of quantum fluctuations. We also analytically derive the condition under which cosmological scaling solutions exist starting from a general Lagrangian including the phantom type scalar field. We apply this method to the case where the dilaton is coupled to non-relativistic dark matter and find that the system tends to become quantum mechanically unstable when a constant coupling is always present. Nevertheless, it is possible to obtain a viable cosmological solution in which the energy density of the dilaton eventually approaches the present value of dark energy provided that the coupling rapidly grows during the transition to the scalar field dominated era. ", "introduction": "One of the most challenging and intriguing problems of cosmology is undoubtedly that of dark energy, the marginally dominant negative pressure component that drives the present acceleration of the Universe (for reviews see e.g. \\cite{varun,paddy}). The identification between dark energy and the energy of the vacuum (cosmological constant), which, in some respect, may look as a ``minimal'' choice, rises at least a couple of embarrassing issues. The main concern is perhaps to explain why the vacuum energy is so small in particle physics units (``fine-tuning'' problem). Moreover, beside being generically ``small'', the cosmological constant happens to be exactly of the value required to become dominant at the present epoch. The latter mysterious circumstance is sometimes referred to as the ``coincidence problem''. These puzzles may be better interpreted by assuming that the energy of the vacuum is, for some unknown cancellation mechanism, exactly zero and by considering in its place a dark energy component with a dynamically variable equation of state. A host of such models have been studied, ranging from quintessence \\cite{quin}, K-essence \\cite{kesse}, braneworlds \\cite{brane}, tachyons \\cite{tachyon}, chaplygin gas \\cite{chap} etc.. Effective scalar fields represent, in this respect, a simple and well motivated choice, since they are omnipresent in supersymmetric field theories and in string/M theory. For instance, a definite prediction of string theory is that the gauge and gravitational couplings are not fixed \\emph{a priori} but, rather, related to the vacuum expectation value of a scalar field, the dilaton. More precisely, at the tree level in the string loop expansion, the vacuum expectation value of the (four dimensional) dilaton $\\phi$ is related to the gauge coupling $\\alpha_{\\rm GUT}$ and to the Planck mass $M_P$ through $e^\\phi \\simeq M_s^2/M_P^2 \\simeq \\alpha_{\\rm GUT}$, where $M_s\\equiv \\sqrt{2/\\alpha '}$ is the string mass and $\\alpha '$ the universal Regge slope parameter of the string. The very existence of a relation between the couplings and the vacuum expectation value (VEV) of a scalar field, although promisingly tasting of ``unification'', threatens with possible violations of the equivalence principle as well as unobserved variations of the coupling constants. It is common wisdom \\cite{TV88} to assume that the dilaton and the other moduli of the theory are given a mass by some non-perturbative mechanism, in such a way that their long range interactions are suppressed and their VEVs frozen at some phenomenologically reasonable value. An alternative possibility is that, at the level of the effective action, the dilaton decouples from the other fields. To be more explicit, let us consider a generic effective string action at lowest order in $\\alpha '$ [we use the convention $(- + + +)$ for the metric] \\begin{equation} \\label{intro} {\\cal S} = \\frac{1}{\\alpha'} \\int \\volt \\, \\left[B_g(\\phi) \\wt R + B_\\phi(\\phi)\\, \\wt{g}^{\\, \\mu\\nu} \\, \\partial_\\mu \\phi \\, \\partial_\\nu \\phi \\, - \\frac{\\alpha'}{4} B_F(\\phi){\\wt F}^2 + \\, \\dots \\, \\right]\\,, \\end{equation} where $\\wt R$ is the scalar curvature and ${\\wt F}$ is the gauge field. Here the dilaton-dependent loop effects as well as the non-perturbative corrections are encoded in the ``coupling functions'' $B_i(\\phi)$, and a tilde denotes the quantities as measured in the ``string frame'', the sigma-model metric minimally coupled to the fundamental strings. In the non--weak coupling region ($e^\\phi~\\gsim~1$), where the loop effects are important, all the (relevant) functions $B_i(\\phi)$ may extremize at some value $\\phi_m$. This possibility has been investigated in Ref. \\cite{DP94} where it has been shown that the cosmological evolution tends to push the dilaton toward the value $\\phi_m$. This is a phenomenologically ``safe'' vacuum where the massless dilaton decouples from the matter fields. By means of a large-N argument, Veneziano has recently suggested \\cite{V01} that the effective couplings $B_i$ are, for the most part, \\emph{induced} by the quantum corrections of the many moduli and gauge bosons of the theory. As a result, these functions of $\\phi$ reach an extremum at infinite \\emph{bare} coupling $e^\\phi \\rightarrow \\infty$ and follow the general behaviour \\begin{equation} \\label{Bidef} B_i(\\phi) = C_i + {\\cal O}(e^{-\\phi}) ~~~~~~~ (e^\\phi \\gg 1). \\end{equation} In this scenario (of the so called ``runaway dilaton''), $\\phi$ gradually decouples from gravity and from the matter fields by evolving towards infinity: $\\phi_m = \\infty$. The moduli of the coefficients $C_i$ are proportional, by factors of order one, to the number $N\\sim 10^2$ of independent degrees of freedom which have been integrated over. What about their signs? It is natural to require a sensible low energy theory e.g. -- in our conventions -- that both $C_g$ and $C_F$ be positive numbers. Moreover, in order for the dilaton $\\phi$ to behave as a ``normal'' scalar field in the limit of $\\phi \\gg 1$, the kinetic coefficient $C_\\phi$ has to be negative. This choice was made in \\cite{GPV,DPV1,DPV2} where the late-time cosmology and phenomenology of this model with a ``well-behaved'' scalar field were studied in some detail. In particular, in \\cite{GPV} it was shown that a residual coupling of the dilaton to some hidden ( --dark matter?) sector of the theory can give rise to a final cosmological attractor with both an accelerated expansion and a constant ratio between dark matter and dilatonic energy densities. Such a type of ``scaling'' behaviour \\cite{wetter,Amendola2}, relaxes the ``coincidence problem'' by interpreting the present energy budget of the Universe (dark matter $\\simeq 1/3$, dark energy $\\simeq 2/3$ of the total) as belonging to a stable attractor configuration. Moreover, these scaling solutions can lead to a viable late-time cosmology \\cite{AGUT} with the acceleration starting earlier ($z>1$) than in usual (uncoupled) dark energy models but still consistent with the recent Type Ia supernovae data \\cite{Amendola3,AGPV}. The opposite choice for the kinetic sign in (\\ref{intro}), $C_\\phi >0$, leads to what in quantum field theory is called a ``ghost'' field or, in a more recent cosmological fashion, a ``phantom'' \\cite{caldwell}. When considered as dark energy candidates, ghosts can have an equation of state parameter $w\\equiv p/\\rho< -1$ which is not ruled out (rather, slightly favoured \\cite{tirthalam}) by current observations. Ghosts/phantoms prove to have a viable cosmological behavior which has been extensively studied in Refs.~\\cite{caldwell,phantom,Sami:2003xv} and also constrained with observations \\cite{sami}. Although intriguing as ``classical -- cosmological'' fields, phantoms are generally plagued by severe UV quantum instabilities\\footnote{The energy of a phantom field is not bounded from below and this makes the vacuum unstable against the production of ghosts + normal (positive energy) fields \\cite{Carroll}. Even when decoupled from the matter fields, ghosts couple to gravitons which mediate vacuum decay processes of the type $vacuum \\rightarrow 2\\, ghosts + 2 \\gamma$ and an unnatural Lorenz invariance breaking -- cut off of $\\sim$ MeV is required to prevent an overproduction of cosmic gamma rays \\cite{Cline}.}, so that the fundamental origin of a $w<-1$ component still represents an interesting challenge for theoreticians. Recently, it has been shown that a scalar field with a wrong sign kinetic term does not necessarily lead to inconsistencies, provided a suitable structure of higher order kinetic terms exists in the effective theory \\cite{Arkani}. The basic mechanism resembles that of a $\\lambda \\vp^4$--theory, where the field acquires a non-zero VEV and the energy of the small fluctuations around the new minimum is bounded from below. By the same token, in a theory with higher order kinetic terms, the field can ``condensate'' at a non-constant background value, $\\dot{\\vp}_0 \\neq 0$, which is perfectly stable at the quantum level. The basic scheme is that of a Lagrangian of the type $p = - \\partial \\vp^2 + \\partial \\vp^4/m^4$ where $\\vp$ is a canonically normalized (ghost) scalar field and $m$ a mass scale. In this case a background value $\\vp_0$ can, at the same time, be quantum mechanically stable ($\\partial \\vp_0^2\\geq m^4/2$) and act as a negative pressure component ($\\partial \\vp_0^2\\leq m^4$). A kinetically driven cosmic acceleration has been proposed both in inflationary \\cite{kinflation,Arkani2} and quintessence \\cite{kesse} contexts. See Refs.~\\cite{ghostrecent} for recent works about ghost condensation. In this paper we highlight the runaway dilaton scenario by exploiting the possibility that string--loop corrections may result in the ``wrong'' kinetic sign for the dilaton at the effective level [i.e. $C_\\phi >0$ in Eq.~(\\ref{Bidef})]. The low energy theory can still be consistent if the dilaton, by a mechanism similar to that in \\cite{Arkani}, ``condensates'' in a ``safe'' vacuum configuration. We also show that a residual coupling of the scalar field to dark matter can lead to a late-time accelerating Universe with the fractional densities of dark matter and dark energy that remain constant in time. In section 2 we sketch the origin of our model from the low-energy limit of string theory. In section 3 we address the stability of quantum fluctuations for a scalar field $\\vp$ whose Lagrangian is a general function of $\\partial \\vp^2$ and of $\\vp$ itself. We also generalize the late-time attractors studied in \\cite{wetter,Amendola2,GPV} and derive the general functional form of the Lagrangians that allow this type of scaling solutions. In section 4 we study the dynamics of the effective ghost/dilaton for various choices of the parameters. It is found that in the absence of the coupling $Q$ between the dilaton and the non-relativistic matter one can get a viable attractor solution which asymptotically approaches a de Sitter phase without violating the stability of the vacuum. We also show that there exist well-behaved cosmological scaling solutions if the coupling $Q$ grows from nearly zero to a constant value during the transition to the dilaton-dominant era. We give conclusions and discussions in the final section. ", "conclusions": "In this paper we have studied a new type of dark energy model based on string theory. Our starting point is the Lagrangian (\\ref{action}) in low-energy effective string theory. We adopted the runaway dilaton scenario in which the dilaton field $\\phi$ is effectively decoupled from gravity as the field evolves toward $\\phi \\to \\infty$. The coefficient of the kinematic term of the dilaton can be negative in the Einstein frame, which means that the dilaton behaves as a phantom-type scalar field. We implemented higher-order derivative terms for the field $\\phi$ which ensure the stability of the system even when the coefficient of $\\dot{\\phi}^2$ is negative. We then considered a general Lagrangian which is generic function of a (canonically defined) scalar field $\\vp$ and $X\\equiv -(\\nabla \\varphi)^2/2$. We first derived the stability condition of quantum fluctuations, see Eq.~(\\ref{xi}). At the classical level the perturbations may be stable if the speed of sound given in Eq.~(\\ref{sound}) is positive, but we obtained a more stringent condition in order to ensure the stability of the vacuum at the quantum level. When the condition (\\ref{xi}) is satisfied, one can avoid the gravitational creation of ghost and photon pairs which was recently raised as a serious problem by Cline et al.\\,\\cite{Cline} for the classical phantom models considered so far \\cite{caldwell,phantom,Sami:2003xv,sami}. In this work we carefully addressed the way to overcome this problem. The stability region is plotted in the phase-space, which restricts the dynamical evolution of the system relative to the normal scalar field with a positive kinematic term. In Sec.\\,\\ref{sec_scaling} we derived a general form of the Lagrangian (\\ref{scap}) for a single field $\\vp$ when the scaling solutions exist. We implemented a coupling between the dilaton and the matter (denoted by $Q$) and obtained the equation of state $w_\\vp$ and the relative energy density $\\Omega_\\vp$ for the field $\\vp$, see Eq.~(\\ref{womega}). This is a general useful expression valid not only for the normal scalar field but also for the ghost-type scalar field. We also found that when the Lagrangian is written as $p=f(X)-V(\\phi)$ the requirement of the scaling solutions uniquely determines the Lagrangian as $p=\\pm X-c_2e^{-\\lambda \\phi}$. This is the system with an exponential potential and a positive/negative kinematic term, which was extensively studied in the literature. We considered the classical dynamics of the system with a negative kinematic term $-X$ and the bell-type potential (\\ref{po}), and reproduced the result that the dilaton evolves toward the potential {\\it maximum} with an equation of state $w_\\vp \\le -1$. However, since this system is unstable at the quantum level, we took into account a term of the type $e^{\\lambda \\vp}X^2$ in order to ensure the stability of quantum fluctuations. When the condition for the stability is satisfied, the equation of state is restricted in the range $w_\\vp \\ge -1$ in general. We found that there exist attractor solutions with $Q=0$ in which $\\Omega_\\vp$ evolves toward 1 and $w_\\vp$ approaches a constant value slightly larger than $-1$. This is different from the scaling solutions discussed in Sec.\\,\\ref{sec_scaling} which keep the relative energy density constant with {\\it nonzero} values of $\\Omega_\\vp$ and $\\Omega_m$. This can be derived by setting $\\rd x/\\rd N=\\rd y/\\rd N=0$ in the evolution equations (\\ref{xeq}) and (\\ref{zeq}). We plotted such examples in Figs.\\,~\\ref{evon} and \\ref{evo}, whose results agree well with analytic estimates. We also checked that the system lies in the quantum stability region shown in Fig.~\\ref{stability} throughout the evolution, thereby providing a viable dark energy model even when the coefficient of the $\\dot{\\vp}^2$ term is negative. When the matter Lagrangian $\\wt{{\\cal S}}_m$ is not dependent on the field $\\phi$ in the string frame, one can show that $Q \\propto (1/B_g) \\rd B_g/\\rd \\phi$ is vanishing as $\\phi \\to \\infty$ for the choice (\\ref{triv}). In this case one can get ideal cosmological solutions as shown in Figs.~\\ref{evon} and \\ref{evo}. If the matter is coupled to the dilaton in the string frame, the coupling $Q$ is not necessarily zero. We analytically found in Sec.\\,\\ref{sec_scaling} that scaling solutions exist when $Q$ is a nonzero constant in the scaling regime. When $Q$ is a constant from the matter-dominated to the dilaton-dominated era, the stability condition is violated for the realistic initial conditions ($x_i^2, z_i^2 \\ll 1$). Therefore it is difficult to obtain a viable cosmological evolution in such a case, see Figs.~\\ref{pspace} and \\ref{phase}. However, if the coupling $Q$ rapidly grows from nearly zero to a constant value as in Eq.~(\\ref{varyQ}), the system can approach the scaling solution without breaking the stability condition, see the case (i) of Fig.~\\ref{varyingq}. In order to get a ``direct'' approach toward $\\Omega_\\vp \\simeq 0.7$ as in the case (iii) of Fig.~\\ref{varyingq}, we require a step-like change of the coupling $Q$ during the transition to the scalar-field dominant era. In our work we carried out the analysis for the simplified Lagrangian (\\ref{ourlag}) in order to understand the basic picture of the system with a phantom-type scalar field. In string theory we have other non-perturbative and loop corrections such as the Gauss-Bonnet curvature invariant \\cite{Cartier,Tsuji,Foffa}. It is certainly of interest to extend our analysis to such a direction." }, "0405/astro-ph0405180_arXiv.txt": { "abstract": "We have completed part of a program to study the X-ray emission properties of old rotation-powered pulsars with XMM-Newton in order to probe and identify the origin of their X-radiation. The X-ray emission from these old pulsars is largely dominated by non-thermal processes. None of the observed spectra required adding a thermal component consisting of either a hot polar cap or surface cooling emission to model the data. The X-ray spectrum of \\PSRB\\, is best described by a single power law of photon-index $\\alpha=1.93^{+0.14}_{-0.12}$. Taking optical data from the VLT FORS1 into account a broken power law model with the break point $E_{break}= 0.67^{+0.18}_{-0.41}$ keV and the photon-index $\\alpha_1= 1.27^{+0.02}_{-0.01}$ and $\\alpha_2 =1.88^{+0.14}_{-0.11}$ for $E < E_{break}$ and $E > E_{break}$, respectively, is found to describe the pulsar's broadband spectrum from the optical to the X-ray band. Three-$\\sigma$ temperature upper limits for possible contributions from a heated polar cap or the whole neutron star surface are $T^\\infty_{pc} < 0.87 \\times 10^6$ K and $T^\\infty_s < 0.48 \\times 10^6$ K, respectively. We also find that the X-ray emission from \\PSRB\\, is pulsed with two peaks per rotation period. The phase separation between the two X-ray peaks is $\\sim 144^\\circ$ (maximum to maximum) which is similar to the pulse peak separation observed in the radio band at 1.4 GHz. The main radio peak and the trailing X-ray peak are almost phase aligned. The fraction of X-ray pulsed photons is $\\sim 30\\%$. A phase-resolved spectral analysis confirms the non-thermal nature of the pulsed emission and finds no spectral variations as a function of pulse phase. Detailed pulse profile simulations using polar gap, outer gap and the two-pole caustic model constrain the pulsar's emission geometry to be that of an almost orthogonal rotator, for which the two-pole caustic model can reproduce the observed doubly peaked X-ray pulse profile. The spectral emission properties observed for \\PSRBB\\, are similar to those of \\PSRB. Its energy spectrum is very well described by a single power law with photon-index $\\alpha=2.5^{+0.9}_{-0.45}$. Three-$\\sigma$ temperature upper limits for thermal contributions from a hot polar cap or from the entire neutron star surface are $T^\\infty_{pc} <1.17 \\times 10^6$ K and $T^\\infty_s < 0.5 \\times 10^6$ K, respectively. There is evidence for pulsed X-ray emission at the $\\sim 97\\%$ confidence level with a pulsed fraction of $49 \\pm 22\\%$. For \\PSRJ\\, which is located $\\sim 1^\\circ$ outside the boundary of the Cygnus-Loop, we report the first detection of X-ray emission. A power law spectrum, or a combination of a thermal and a power law spectrum all yield acceptable descriptions of its X-ray spectrum. No X-ray pulses are detected from \\PSRJ\\, and the sensitivity is low - the $2\\sigma$ pulsed fraction upper limit is $57\\%$ assuming a sinusoidal pulse profile. ", "introduction": "} In the past decade, many advances have been made in the study of rotation powered pulsars thanks to X- and $\\gamma$-ray observations with a series of high energy astrophysical space missions (for reviews see e.g.~Becker \\& Tr\\\"umper 1997; Becker \\& Pavlov 2001; and Kanbach 2001). However, many outstanding scientific issues still remain to be resolved. Understanding the high energy emission processes of old rotation-powered pulsars is one of them. Although one could use the instruments aboard ROSAT, BeppoSAX and ASCA to disentangle the thermal and non-thermal contributions of the young and cooling neutron stars, the sensitivity of these instruments was not sufficient to study the emission properties of the old field pulsars. Old rotation-powered field pulsars are of particular interest for the study of particle acceleration and high energy radiation processes on the neutron star surface and in the neutron star magnetosphere. This is because their ages are intermediate between that of the young cooling neutron stars, whose surface may produce copious thermal X-ray photons, and those very old millisecond pulsars, in which non-thermal magnetospheric X-ray production mechanisms are believed to dominate (see e.g. Becker \\& Tr\\\"umper 1999, Becker et al.~2003, Webb et al.~2004). The old field pulsars aid in answering questions such as \"How do the emission properties of the younger pulsars like Geminga, PSR 0656+14 and PSR 1055-52 change as they age from $\\sim 10^5$ to $10^7$ years\"? Will the thermal emission simply fade away due to cooling with increasing age or will the star be kept hot (at about $0.5-1 \\times 10^5$ K) over millions of years due to energy dissipation by processes such as internal frictional heating ($\\dot{E}_{diss} \\sim 10^{28}-10^{30}$ erg/s) and crust cracking, as proposed by Alpar et al.~(1984; 1998) and Ruderman (1998) or by vortex creeping and pinning models (Shibazaki \\& Lamb 1989; Hirano and Shibazaki, 1997)? What happens to the non-thermal, hard-tail emission seen in the X-ray spectra of the middle aged field pulsars? Will this emission be the dominant source or will this component also decay with time and will only thermal emission from the hot and heated polar-caps remain? Establishing whether thermal polar-cap emission is present, or not, is of importance in confronting and comparing data with the many magnetospheric emission models which predict hot polar caps (e.g.~Arons \\& Scharlemann 1979; Sturner \\& Dermer 1995, Zhang \\& Harding 2000; Harding \\& Muslimov 2001; 2002; 2003). The heated polar cap is a consequence of pair-creation by the gap discharge, after which a significant amount of highly energetic charged particles is expected to stream back to the neutron star, heating the surface to a few million degrees. Furthermore, many magnetospheric emission models {\\em require} hot-polar caps in order to lessen the work-function of electrons and positrons in the surface down to $\\sim 100$ eV and thus make it possible to pull out enough electrons for the vacuum gap discharge/break-down (cf.~Michel 1991). These models fail if there are no hot-spots. On the other side, Sturner \\& Dermer (1995) propose that high energy gamma-ray photons are created in pulsar magnetospheres by inverse Compton scattering of relativistic electrons and thermal photons. If this model is correct, the absence of hot polar-caps (i.e.~the missing bath of thermal photons) would imply that old -- but close-by and still powerful -- pulsars are not $\\gamma$-ray emitters, in agreement with the current observations. There are further questions that involve the old pulsar's energy output. The observed luminosity is but a small fraction of the total energy available due to the rotation of the star. Where is the bulk of the pulsar's spin-down power going to? Are there pulsar-wind nebulae around these systems not previously detected due to the sensitivity limitations imposed by previous satellites? Up to now, only three old, non-recycled, pulsars have been detected and all the detections were close to the sensitivity limits of the instruments, thus strongly limiting the ability to explore the physical emission processes at work in these neutron stars. The detected pulsars are PSR B1929+10 (Helfand 1983; Yancopoulos, Hamilton \\& Helfand 1994), B0950+08 (Seward \\& Wang 1988; Manning \\& Willmore 1994; Saito 1998) and B0823+26 (Sun et al.~1993). All have a spin-down age of {\\large $\\tau$} $\\sim 0.2-3 \\times 10^7$ years; magnetic fields of about $B_\\perp\\sim 10^{11}-10^{12}$G; a close distance of $d \\sim 0.12 - 0.38$ kpc and a very small absorption column of $N_H\\sim 10^{18}-10^{20}\\; \\mbox{cm}^{-2}$ (see e.g.~Hobbs et al.~2003). Spectral and temporal information was only available for the brightest of these, PSR 1929+10. Its X-ray pulse profile is very broad with a single pulse stretching across the entire phase cycle. The fraction of pulsed photons is $\\sim 30\\%$ (Yancopoulos, Hamilton \\& Helfand 1994). Its X-ray spectrum, observed with ROSAT, could be equally well fit with a power law (photon-index $\\alpha\\sim 2$) and a black-body spectrum (thermal polar-cap emission, $T\\sim 3.2\\times 10^6$ K, $R_{bb} \\sim 20-50$\\,m), leaving open the real nature of its emission (Becker \\& Tr\\\"umper 1997). Further evidence for a non-thermal nature of the pulsar's X-ray emission was found by Saito (1998) based on ASCA data and more recently by Wozna et al.~(2003) in a joint analysis of archival ROSAT and ASCA data. Both, PSR 1929+10 and \\PSRB\\, are detected at optical wavebands (Pavlov et al.~1996; Zharikov et al.~2002; Mignani et al.~2003). Making use of the large collection area, as well as the excellent timing, spatial and energy resolution provided by XMM-Newton, the three pulsars \\PSRB, \\PSRBB\\, and \\PSRJ\\, which all belong to the sub-class of old rotation-driven field pulsars, were observed as part of the guaranteed time and AO1 guest observer program. We list the radio properties of these pulsars in Table~\\ref{t:radio} and make the comments that follow. Based on its spin-down age, \\PSRB\\ is the oldest field pulsar among the more than 50 rotation-driven pulsars detected in X-rays (see e.g.~Table 3 of Becker \\& Aschenbach 2002 for a recent list of X-ray detected rotation-powered pulsars). The pulsar's radio dispersion measure is among the smallest of all known radio pulsars and the dispersion-measure-based distance is is in good agreement with the $(262 \\pm 5)$pc distance deduced from a parallax measurement (Brisken et al.~2002). The low column density implies that it should be feasible to detect soft (below 0.5 keV) X-ray emission from \\PSRB. X-rays from \\PSRB\\, were first detected by Maning \\& Willmore (1994) in a $\\sim 9\\,\\mbox{ksec}$ ROSAT PSPC observation ($\\sim 55$ source counts). These authors suggested that the emission either came from a $T\\sim (2.1\\pm 0.2) \\times 10^6$ K hot polar cap of $R_{bb}\\sim 20$m or arose from synchrotron or curvature radiation with a power law spectrum of photon-index $\\alpha=0.9^{+2.2}_{-1.3}$. Interpreting the soft X-ray emission entirely as arising from the neutron star's surface, Becker (1994) computed a $3\\sigma$ temperature upper limit of $T_s^\\infty < 0.23 \\times 10^6$ K assuming a 1.4 M$_\\odot$ neutron star with a medium stiff equation of state. \\PSRBB\\, has an inferred magnetic field strength that is the highest among the X-ray-detected old field pulsars. X-ray emission from \\PSRBB\\, was first discovered using ROSAT (Sun et al.~1993). The ROSAT PSPC count rates were $0.0015 \\pm 0.0004$ cts/s and $0.0009 \\pm 0.0003$ cts/s for the $(0.8-2.0$ keV) and $(0.5-2.0$ keV) bands, respectively. The small number of detected source counts did not allow one to identify the emission process but the data were in agreement with the hypothesis that the X-rays were emitted from a $\\sim 1.8\\times 10^6$ K hot thermal polar cap of size $R_{bb}\\sim 140$m. If the soft X-rays were assumed to arise entirely from the neutron star's surface, a $3\\sigma$ temperature upper limit is $T_s^\\infty <0.34 \\times 10^6$ K assuming a 1.4 M$_\\odot$ neutron star with a medium stiff equation of state (Becker 1994). \\PSRJ\\, (Thorsett et al.~1994; Camilo \\& Nice 1995; Ray et al.~1996) is another important representative of the group of old, but non-recycled, pulsars. Its rotation period is the shortest among the old rotation-powered field pulsars. Compared to \\PSRB\\ and \\PSRBB , its spin-down energy is two orders of magnitude higher whereas its inferred magnetic field strength is similar to that computed for \\PSRB. The pulsar's spin-down age makes it intermediate between the older pulsars \\PSRB\\ \\& \\PSRBB\\ and the cooling neutron stars, which all have a spin-down age of some hundred thousand years (see Table 3 of Becker \\& Aschenbach 2002). Its radio dispersion-measure inferred distance suggest a medium to low neutral hydrogen column density. The moderate column density is an advantage in searching for X-ray emission from this pulsar, which has never been previously observed by any high energy mission and was not known to emit X-rays. \\PSRJ\\ is located about $1^\\circ$ outside the boundary of the Cygnus-Loop supernova remnant but an association would require that the pulsar's true age is rather different compared with its spin-down age and that it was born with a very high velocity. In \\S2 we describe the XMM-Newton observations of \\PSRB, \\PSRBB\\, and \\PSRJ\\, and provide the details of the data processing and data filtering. The results of the spectral and timing analysis are given in \\S3 - \\S5. We provide a summary and concluding discussion in \\S6. ", "conclusions": "} We have investigated the X-ray emission properties of three old, but non-recycled, rotation-powered pulsars in order to probe and identify the origin of their X-radiation. These pulsars, being intermediate in age between the young cooling neutron stars and the old recycled millisecond pulsars, are of special interest as they provide important information for understanding the X-ray emission properties of rotation-powered pulsars as a class. The selected targets provide a valuable snapshot at ages {\\Large $\\tau$}$ = 1.2 \\times 10^6$ years (\\PSRJ), {\\Large $\\tau$}$ = 4.89 \\times 10^6$ yrs (\\PSRBB) and {\\Large $\\tau$}$ = 1.74 \\times 10^7$ yrs (\\PSRB) and allows one to add to the current picture of pulsar X-ray emission properties beyond the younger (spin-down ages between $1-6 \\times 10^5$ years) class of cooling neutron stars. For \\PSRB, which is the oldest among the three pulsars investigated, any hint of emission from the cooling stellar surface has faded to below what might be detected in the XMM-Newton observation. The $3\\sigma$ surface temperature upper limit of $T_s^\\infty < 480\\,000$ K is well above temperatures predicted by current models of neutron star thermal evolution even if strong frictional heating of superfluid ${}^1\\mbox{S}_0$-neutrons in the outer neutron star crust is considered (Umeda et al.~1993; Yakovlev et al.~2002). The same is true for \\PSRBB\\, where $T_s^\\infty < 500\\,000$ K and for \\PSRJ\\, for which $T_s^\\infty < 627\\,000$ K. However, standard neutron star cooling models neglect the influence of a strong magnetic field on the neutron star's thermal evolution. As the heat transport in neutron stars is mainly due to electrons the presence of a magnetic field is supposed to reduce the thermal conductivity perpendicular to the magnetic field direction. The consequences are an anisotropic temperature distribution on the neutron star surface (Geppert, K\\\"uker \\& Page 2004) and a reduced cooling rate so that magnetic cooling curves may deviate significantly from the zero-field case after $10^{5-6}$ years (Tsuruta 1998). The surface temperature upper limits of the old, non-recycled, pulsars, even if they are above what standard cooling models predict, may still provide interesting constraints for those thermal evolution models which take the neutron star's magnetic field into account. As far as emission from a thermal polar cap is concerned, it is very interesting that there is no clear evidence for the presence of this component in any of the energy spectra. Yet, emission from a heated polar cap is present in the spectra from both the younger cooling neutron stars and the older, recycled, millisecond pulsars (cf.~Becker \\& Pavlov 2001) as long as the neutron star is active as a pulsar. Harding \\& Muslimov (2001; 2002) predicted in the framework of their revised space-charge-limited flow model that polar cap heating, as a fraction of the spin-down luminosity, increases with pulsar age and should be most efficient for pulsars of spin-down age {\\Large $\\tau$}$\\sim 10^7$ yrs, if they are in fact producing pairs from curvature radiation photons. According to these models, however, B0950+08 and B0823+26 cannot produce pairs from curvature radiation (CR) of primary electrons since they both lie below the CR pair death line in the $P$-$\\dot{P}$ diagram of radio pulsars (i.e.~the primary electrons cannot accelerate to the energies required to produce CR pairs). Both of these pulsars can however produce pairs from inverse Compton scattered (ICS) photons, which provide much lower PC heating than do CR-produced positrons resulting in predicted luminosities of $L_{+}^{ICS} \\simeq 10^{28}\\,\\rm erg\\,s^{-1}$ for B0950+08 and $L_{+}^{ICS} \\simeq 6 \\times 10^{27}\\, \\rm erg\\,s^{-1}$ for B0823+26. Both of these values are well below the luminosities that we have observed for these sources, and also below the upper limits for emission from a heated polar cap which are $L_{pc} < 8.4 \\times 10^{28}\\,\\rm erg\\,s^{-1}$ for B0950+08 and $L_{pc} < 1.3 \\times 10^{29}\\,\\rm erg\\,s^{-1}$ for B0823+26. The results are thus consistent with a non-thermal, not a polar cap heating, origin for the emission from these two pulsars. J0243+2740 lies well above the CR pair death line of Harding \\& Muslimov (2002), and thus is expected to have a much higher level of polar cap heating from CR produced positrons with a luminosity predicted to be $L_{+}^{CR} \\simeq 10^{31}\\,\\rm erg\\,s^{-1}$. This value is near but below our observed luminosity of $L_{obs} =2.8 \\times 10^{31}\\,\\rm erg\\,s^{-1}$, implying that a significant part of the observed luminosity could come from polar cap heating, depending on how directly we are viewing the polar cap. Thus the luminosities detected from all three pulsars are consistent with the predicted level of polar cap heating. The geometry of \\PSRB\\, and \\PSRBB\\, has been investigated recently by fitting the classical rotating vector model to high-quality polarization data taken with Arecibo Observatory at 1.4 GHz (Everett \\& Weissberg 2001). These authors favor the interpretation that both pulsars are almost orthogonal rotators, but Narayan \\& Vivekanand (1982), Lyne \\& Manchester (1988), Blaskiewicz, Cordes \\& Wassermann (1991), Rankin (1993a; 1993b) and von Hoensbroech \\& Xilouris (1997) in previous observations came to the conclusion that the emission geometry of \\PSRB\\, is that of an almost aligned rotator. Figure \\ref{PSRB_geometry} shows the geometry of the two scenarios for \\PSRB\\, with inclination and impact angles taken from Everett \\& Weissberg (2001) and references therein. Although rotating vector model fits are easily perturbed by systematic effects in polarized position angles and reported uncertainties often underestimated the actual errors, we find that the double peaked X-ray pulse profile of \\PSRB\\, with the peak separation $\\delp \\simeq 0.4$ strongly supports the nearly orthogonal rotator. In the aligned rotator geometry, with an inclination angle of $\\sim 170^\\circ$ and impact angle of $\\sim 5^\\circ$ (see Figure \\ref{PSRB_geometry}a), the polar cap model (Ruderman \\& Sutherland 1975; Daugherty \\& Harding 1982) predicts a single-peaked profile. Double peaked profiles (as those in Daugherty \\& Harding 1996, and Dyks \\& Rudak 2002) can be observed only when our line of sight crosses the polar gap. This would require the gap to be located at least 15 stellar radii above the surface, or the surface conal beam would have to be 4 times wider than the polar cap beam. Moreover, an improbably fine tuning of model parameters would be required to reproduce the large peak separation. In the outer magnetosphere scenarios such as the outer gap model (Cheng, Ho \\& Ruderman 1986; Romani \\& Yadigaroglu 1995; Cheng, Ruderman \\& Zhang 2000) or the two-pole caustic model (Dyks \\& Rudak 2003), the non-thermal X-rays are emitted in a fan beam. For the nearly aligned geometry, however, the outer gap model predicts no high-energy radiation (see top panel in Figure 6 in Cheng, Ruderman \\& Zhang 2000) whereas the two-pole caustic model predicts single-peaked lightcurves (see Figure 2b in Dyks \\& Rudak (2003). Both the outer gap and two-pole caustic models can reproduce the observed profile in the nearly orthogonal scenario. Figure \\ref{tpc} presents the radiation pattern ({\\it top}) and the pulse profile ({\\it bottom}) calculated for the two-pole caustic model with the dipole inclination $\\alpha = 105^\\circ$ and and the viewing angle $\\alpha+\\beta = 127^\\circ \\equiv \\zeta$, as derived by Everett \\& Weisberg (2000). Each peak arises due to the caustic effects on the trailing side of the open field line region associated with each magnetic pole (see Figure \\ref{tpc}a). The modeled peak separation ($\\delp \\simeq 0.43$) is in good agreement with the observed one ($\\sim 0.4$). The {\\it relative} widths and heights of these two peaks also resemble the observed ones, however, they are more model-dependent than $\\delp$. Our three-dimensional simulations show that for the parameters given above, the outer gap model also predicts a double-peaked profile, with $\\delp \\le 0.3$ which is marginally consistent with the data. The outer gap model can more closely reproduce the large peak separation for viewing angles closer to the rotational equator ($\\zeta \\sim 100^\\circ$). However, according to outer gap models, only younger pulsars can sustain a gap in their magnetospheres and produce non-thermal high-energy emission. B0950+08, B0823+26 and J0243+2740 all lie below the original outer gap death lines for production of high-energy emission (Chen \\& Ruderman 1993), indicating that they do not produce high-energy emission from outer gaps. More recently, the outer gap death lines have been revised to include pulsar inclination and thermal emission from cooling and heated polar caps (Zhang et al.~2004). All three pulsars lie above at least one of the revised outer gap death lines computed by Zhang et al.~(2004), so that outer gap emission may not be ruled out. Neither the two-pole caustic nor the outer gap model can explain the relative locations of the X-ray and radio peaks. Both models predict that the main (i.e., the strongest) radio peak should precede the leading X-ray peak in phase roughly by $\\sim 0.1$. This is the phase at which our line of sight approaches most closely one of the magnetic poles (Figure \\ref{tpc}a). The standard polar cap model in the nearly orthogonal geometry cannot explain the X-ray lightcurve. Because the closest approach to a magnetic pole occurs near the leading peak, the model predicts that this peak should be stronger and more spiky than the trailing peak. In the case of \\PSRBB\\, multiple emission components are seen in the radio pulse profile while the statistics of the available XMM-Newton data is not sufficient to better resolve the X-ray pulse profile than to a single broad peak. Non-thermal X-radiation processes implied by the outer magnetosphere interpretation are in agreement with the non-thermal spectra which dominate the emission from all three pulsars. This is most evident in the energy spectrum of \\PSRB\\, but appears to be the case for both \\PSRBB\\, and \\PSRJ\\, as well. However, the predicted level of polar cap heating for \\PSRJ\\ implies that a significant part of its emission may be thermal. The optical emission from \\PSRB\\, has been recently observed with the VLT FORS1 (Zharikov et al.~2003). Fitting the optical data simultaneously with the X-ray spectrum suggests a broadband spectrum which can be described by a broken power law, strongly suggesting that the radiation from \\PSRB\\, is dominated by non-thermal emission from the optical to the X-ray band. Even more, taking \\PSRB, \\PSRBB\\, and \\PSRJ\\, as representative would imply that the X-ray emission from old, non-recycled, rotation-driven pulsars is dominated by non-thermal radiation as has been concluded by Becker \\& Tr\\\"umper (1997) based on their tight soft-X-ray luminosity vs.~spin-down energy conversion fits of ROSAT detected rotation-powered pulsars. The emission properties observed from PSR B1929+10 (cf.~\\S \\ref{intro}) are not in disagreement with this conclusion. XMM-Newton observations of this pulsar, which have taken place in November 2003 and April 2004, are expected to further constrain this conclusion." }, "0405/astro-ph0405463_arXiv.txt": { "abstract": "VLT images in $BVI$ are used to identify the ionizing source centered on Sersic 13, the largest HII region of the giant nearby galaxy NGC~5128 with $log L_{H\\alpha}=39.6$ erg/s. This ionizing source turns out to be a close pair of bright and blue star cluster candidates. Spectroscopy obtained with the Magellan~I telescope confirms that these are massive young clusters physically associated with the giant HII region Sersic 13. The spectra of both clusters show prominent Wolf-Rayet type emission features, and prominent lines of HI and HeI, indicative of a very young age ($t \\approx few\\times 10^6$ yr). Their luminosities make each of them at least as luminous as the massive young cluster R136 in 30 Doradus in the LMC, and their individual masses are estimated to be $1-7.5\\times 10^5$ M$_{\\odot}$. In addition, the projected separation of the cluster pair is $42$ pc. The measured velocity difference between the clusters is small, $\\Delta V = 49\\pm 21$ km/s, and within 2$\\sigma$ of the expected orbital velocity $V_{orb}=5-12$ km/s if they are bound. Dynamical models predict that binary clusters with these properties would merge in a short timescale of a few orbital periods ($P = 20-50\\times 10^6$ yr). The discovery of this binary cluster suggests that mergers of young massive clusters could lead to the formation of the most massive globular clusters such as $\\omega$Cen in our Galaxy and $G1$ in M31. Alternatively, if they are not gravitationally bound, these objects would individually evolve into two normal globular clusters. ", "introduction": "The young populous clusters were defined as a special class by Hodge (1961), who studied blue compact clusters in the Large Magellanic Cloud (LMC). Kennicutt \\& Chu (1988) concluded that populous blue clusters may form in the centers of giant HII regions, such as the central massive cluster R136 in the giant HII region 30Doradus in the LMC. Due to its proximity, R136 is the best studied representative of this class. The properties of the populous blue clusters have been discussed extensively in the literature (Ma\\'{\\i}z-Apellan\\'{\\i}z 2002 and references therein). They are massive, $M>10^5M_{\\odot}$, luminous, $M_V<-10$, young, $t<10^7$ yr, and are often embedded in giant HII regions. For an on-going spectroscopic survey of stellar clusters in the inner regions of the nearby giant elliptical galaxy NGC~5128, we selected candidates based on BVI images obtained with the VLT. In particular, we targeted candidates apparently centered on the brightest HII regions of this galaxy, which have been studied by several authors in the past (see Phillips 1981). The largest of these HII regions is located at the northern edge of the dust disk of this galaxy, at $RA(2000) = 13:25:27.5$, $DEC(2000)=-43:00:11$ with a size of about $0.7\\times 1.0$ kpc. This largest NGC~5128 HII region has been named Nr.\\ 13 by Sersic (1969), Nr.\\ 13 by Moellenhoff (1981), Nr.\\ 10 by Graham (1979), and Nr.\\ 34 by Dufour et al.\\ (1979). Even though the HII region Sersic 13 has been observed before, the nature of its exciting central source was not realized. Phillips (1981) and Rosa \\& D'Odorico (1986) noticed that the WR features present were indicative of the presence of bright young stars. Moellenhoff (1979) reports the presence of two stellar knots in the center of this region based on narrow-band photography, and speculates that they are two bright O-type stars. We extend their work here by studying the photometric and spectroscopic properties of the components of this exciting source, showing that it is composed of 2 very massive young populous clusters, separated by 2.25 arcsec in the N-S direction. In order to avoid adding yet another notation, we call the clusters Sersic 13-N and Sersic 13-S. Typical globular clusters are made of stars formed together from one large parent cloud. Their constituent stars share all the same chemical composition and age. The most massive globular clusters, such as $\\omega$Cen in the Milky Way, and $G1$ in M31, are an exception, and their formation may have been different from the bulk of the population. Following the confirmation that the stellar population in these clusters is composite (Hilker \\& Richtler 2000, Pancino et al. 2000, Meylan et al. 2001), three different formation theories are considered for such massive clusters. One is the merger of binary clusters (Icke \\& Alcaino 1988, Sugimoto \\& Makino 1989, van den Bergh 1996), the second is the stripping of nucleated dwarf galaxies (Freeman 1993, Meylan 2002), and the third involves self-enrichment (e.g.\\ Morgan \\& Lake 1989, Parmentier et al.~1999). While the young binary star clusters appear to be rather frequent in smaller galaxies, like the Magellanic Clouds (e.g.~Bhatia \\& Hatzidimitriou 1988, Bhatia et al.\\ 1991, Dieball, M\\\"uller \\& Grebel 2002), no binary globular clusters have been found. Either they do not form because globular clusters generally trace kinematically hot populations, where the possibility of capture is rather small, or they do not last long because close binary clusters merge rapidly after they form. There are no known binary globular clusters in the Milky Way, which contains 150 globular clusters in total. Perhaps binary clusters are just rare, and one should search in a larger population. The peculiar giant elliptical galaxy NGC~5128 (Centaurus A) holds a large globular cluster system, $>$10 times that of the Milky Way (Kissler-Patig 1997). In this paper we report, for the first time, the identification of a massive young binary star cluster pair in NGC~5128, using images obtained with the ESO VLT, and spectroscopy acquired with the Magellan I telescope. The binary cluster subject of this study is a very interesting object, because, as we discuss below, it may lead to the formation of a very massive globular cluster like $\\omega$Cen or G1. The paper is organized as follows. Section 2 gives the details of the photometric and spectroscopic observations and reductions, and determine the physical properties of our targets. In Section 3 we compare these targets with the young massive cluster R136 of 30Doradus in the LMC. Section 4 discusses the possibility that the most massive object could be a young binary globular cluster and Section 5 summarizes the results of this work. ", "conclusions": "We identify a candidate binary massive cluster in the inner region of NGC~5128. This pair of clusters is centered on the largest HII region of the galaxy, Sersic 13, and we baptize this as the most exciting binary cluster of this galaxy. The components have been classified as young massive clusters on the basis of their sizes, magnitudes and colors, and the reliability of the identification has been confirmed spectroscopically. Both of them have Wolf-Rayet type spectra, and are at least as luminous as R136 in the LMC. The measured radial velocity difference ($\\Delta V = 49$ km/s), and projected separation ($\\Delta s = 42$ pc), are consistent with a binary object within the errors. Kennicutt \\& Chu (1988) suggested that giant HII regions such as 30Doradus in the LMC can be the birth places of massive young globular clusters. In this paper we extend this concept, because the discovery of this binary cluster suggests that, at least in some cases, mergers of young massive clusters could lead to the formation of the most massive globular clusters such as $\\omega$Cen in our Galaxy and $G1$ in M31. Alternatively, if they are not gravitationally bound, these objects would individually evolve into two normal globular clusters. Their evolution depends on their estimated masses $>10^5 M_{\\odot}$, which are very uncertain. Dynamical masses based on integrated high-dispersion spectroscopy are needed to constrain the masses of these clusters in NGC~5128." }, "0405/astro-ph0405149_arXiv.txt": { "abstract": "The motion of sound waves propagating in the perfect fluid with inhomogeneous background flow is effectively described as a massless scalar field on a curved space-time. This effective geometry is characterized by the acoustic metric, which depends on the background flow, and null geodesics on the geometry express the acoustic causal structure. Therefore by the effective geometry we can easily study the causality on the flows. In this paper, we consider a spherically symmetric, relativistic outflow and present the maximal causally connected region for a super-sonic flow. When Lorentz factor of the radial velocity of the flow is constant or obeys power-law with respect to the radial coordinate $r$, we can solve it analytically. As a result we show that in the constant case the maximum angle is proportional to inverse of Lorentz factor and logarithmically increases with respect to $r$, in contrast, accelerative expansions in power-law case make this angle bounded. ", "introduction": "Recently a highly polarized prompt $\\gamma$-ray emission was reported by RHESSI observation of the $\\gamma$-ray burst (GRB) GRB021206 \\cite{Coburn}. This result has inspired many discussions; some authors contravene this analysis of the observation \\cite{Rutledge} and others suggested various models to produce such a high polarization \\cite{Waxman,Granot,Lyutikov,Lyutikov2,Nakar}. For instance, in the electromagnetic models considering Poynting flux dominated flows, polarization arises from a uniform, large scale magnetic field \\cite{Lyutikov,Lyutikov2}. On the other hand, in the hydrodynamic models (e.g., the fireball model \\cite{Piran}) it was suggested that polarization arises even from a random magnetic field \\cite{Waxman}. Besides polarized emissions have been detected also in the GRB afterglows and there are many discussions whether ordered magnetic fields exist or not \\cite{Covino,Gruzinov,Matsumiya,Sari,Medvedev,Rossi}. If polarizied emission needs large scale coherence of magnetic fields, acoustic causality inside the outflow is important to maintain the coherence. The acoustic causality of a magnetic-dominated outflow was investigated in electromagnetic models of GRBs \\cite{Lyutikov,Lyutikov2}. However these analyses were not enough because the effect on the waves propagating in inhomogeneous flows was neglected. For studying such propagations in inhomogeneous flows it is very useful to utilize the effective geometry. The reason is as follows. Sound wave propagation in an inhomogeneous flow of a perfect fluid or that of electromagnetic waves in an inhomogeneous medium are both equivalent to propagating massless scalar fields on an effective curved space-time. These effective geometries are determined by background quantities including the real geometry of spacetime. This idea has been applied to various systems: fluids (ordinary \\cite{Unruh} or super \\cite{Volovik}), dielectrics \\cite{Reznik}, non-linear electromagnetism \\cite{Novello}, and Bose-Einstein condensates \\cite{Garay,Liberati}. Furthermore effective geometries permit us to study non-gravitational systems with methods and ideas of general relativity, such as geodesics, light cone, black holes, ergo-spheres, Hawking radiation, and so on \\cite{Unruh,Visser,Volovik,Garay,Liberati,Schutzhold,Leonhardt,Sakagami,Barcelo,Novello,Reznik}. Particular applications in astrophysics are non-linearmagnetism near a magneter and hydrodynamic accretion-flow onto a black hole \\cite{Bergliaffa,Moncrief}. In this paper, we apply the effective geometry to the causality analysis of relativistic outflows from high-energy sources, which occurs, for example, with GRBs. In Sec.~\\ref{sec:formalism} we review the acoustic metric associated with the effective geometry for sound propagation in a relativistic fluid \\cite{Moncrief,Bilic}. In Sec.~\\ref{sec:analysis}, by the effective geometry we shall analyse the causal structure of a spherical outflow in case that radial Lorentz factor is constant or power-law with respect to the radial coordinate. ", "conclusions": "} We have studied the acoustic causality in relativistic spherical outflows with constant and power-law Lorentz factor case by means of the effective geometry. The effective metric made the analysis of a sound propagation on inhomogeneous background flows very easy. When the Lorentz factor of the outflow is constant, the maximum polar angle of the causally connected region is proportional to the inverse of the Lorentz factor of the outflow and the logarithm of the radius. For a constant or decelerative expansion the maximum angle increases as time without a bound. However for accelerated expansion this angle will reach a constant value at late time and this means the region outside this angle is never causally connected. These results are very similar to the expanding universe in the cosmology. If a shell has a finite thickness sound waves may reach the edge of the shell. For a given thickness we can consider only the region between both edges of the shell. Therefore the previous trajectory with the maximal polar angle is outside the shell. In this case the causally connected region will be narrower than the case without edges. In this paper, we dealt with simple examples. However if a model of GRBs is given and background flow is determined we can investigate causality in a similar way. We expect that even in MHD or fluids with vorticity these results are basically unchanged \\cite{Bergliaffa2}." }, "0405/astro-ph0405239_arXiv.txt": { "abstract": "The $z=3.02$ quasar SDSS J095253.83+011421.9 exhibits broad metal-line emission (\\civ\\ FWHM$\\simeq$9000 \\kms), but broad \\lya\\ emission is not present. Instead, only a narrow \\lya\\ line is observed (FWHM$\\simeq$1140 \\kms). The large \\civ/\\lya\\ ratio in the broad-line region (BLR) emission from this object can be matched most closely by a BLR dominated by gas at very high densities ($10^{15}$\\,cm$^{-3}$), which suppresses the \\lya\\ emission, and illuminated by an incident power-law extending to $\\sim$200\\,\\micron, which yields increased emission from purely collisionally excited coolant lines (such as \\civ, \\Nv\\ and \\ovi) but not from recombination lines like \\lya. However, the strong \\Ciii\\ emission predicted by this model is not observed, and the observed broad \\ciii\\ emission must come from a lower-density BLR component and should be accompanied by broad \\lya\\ emission which is not observed. The least unlikely explanation for this spectrum seems to be that any intrinsic broad \\lya\\ emission is removed by smooth \\Nv\\ absorption in the red wing of the \\lya\\ emission line and by smooth \\lya\\ absorption in the blue wing of the \\lya\\ emission line. This postulated smooth absorption would be in addition to the strong, associated, narrow absorption seen in numerous ions. Smooth absorption in \\lya, \\Nv\\ and \\ovi\\ but not in \\civ\\ would be unusual, but not impossible, although it is suspicious that the postulated absorption must almost exactly cancel the postulated intrinsic broad emission. We conclude that the spectrum of \\nlya\\ appears unique (among $\\simeq$3600 SDSS spectra of quasars at $z>2.12$) because of some {\\em combination} of unusual parameters, and we discuss possible observations to determine the combination of circumstances responsible for the spectrum. ", "introduction": "\\label{INTRO} One of the goals of the Sloan Digital Sky Survey \\markcite{yor00}(SDSS; {York} {et~al.} 2000) is to obtain spectra for $\\sim$10$^5$ quasars, in addition to the $\\sim10^6$ galaxies which comprise the bulk of the spectroscopic targets. From astrometrically calibrated drift-scanned imaging data \\markcite{gun98,sdss153}({Gunn} {et~al.} 1998; {Pier} {et~al.} 2003) on the SDSS $ugriz$ AB asinh magnitude system \\markcite{fuk96,sdss26,sdss82,sdss85,sdss105}({Fukugita} {et~al.} 1996; {Lupton}, {Gunn}, \\& {Szalay} 1999; {Hogg} {et~al.} 2001; {Stoughton} {et~al.} 2002; {Smith} {et~al.} 2002), quasar candidates are selected primarily using color criteria designed to target objects whose broad-band colors are different from those of normal stars and galaxies \\markcite{sdssqtarget}({Richards} {et~al.} 2002). The First Data Release of the SDSS (DR1; \\markcite{dr1}{Abazajian} {et~al.} 2003) includes fluxed, wavelength-calibrated spectra of $\\sim$17,000 quasars \\markcite{dr1q}({Schneider} {et~al.} 2003). This enormous sample includes some quasars with unusual properties. One example is SDSS J095253.83+011421.9 (hereafter SDSS J0952+0114), whose spectrum shows both broad and narrow metal-line emission but only narrow \\lya\\ emission (Fig. \\ref{f_spec}). In this paper we investigate possible explanations for the missing broad \\lya\\ in \\nlya, namely: dust in the broad-line gas (\\S\\,\\ref{DUST}), anisotropic \\lya\\ emission (\\S\\,\\ref{RAD}), broad-line gas with unusual physical properties such that \\lya\\ is intrinsically weak (\\S\\,\\ref{PAR}), an unusual spectrum incident on the broad-line region (\\S\\,\\ref{IR}), or an absorption effect wherein \\Nv\\ and \\lya\\ absorption remove the red and blue wings of broad \\lya, respectively (\\S\\,\\ref{ABS}). We summarize our conclusions in \\S\\,\\ref{CON}. ", "conclusions": "\\label{CON} We have discussed the unusual spectrum of \\nlya, which exhibits broad metal-line emission but only narrow \\lya\\ emission. We argue that this unusual spectrum cannot be {\\em solely} due to dust extinction in the BLR, to anisotropic emission of \\lya, or to unusual physical conditions in the BLR. Most but not all of the spectrum's properties are explainable as emission from a BLR of predominantly high density gas ($n_H\\sim10^{15}$ cm$^{-3}$), which suppresses \\lya, illuminated by an incident power-law continuum extending to $\\geq$200\\,\\micron, which increases the collisionally excited metal-line emission. % However, the BLR in \\nlya\\ cannot consist exclusively of high density gas because the observed broad \\ciii\\ emission would be collisionally deexcited at densities higher than $10^{12}$ cm$^{-3}$. Although it is not entirely satisfactory, the most plausible explanation we have found for the apparent lack of broad \\lya\\ is that it is due to smooth absorption by \\Nv\\ in the red wing of \\lya\\ and by \\lya\\ in the blue wing of \\lya. Such absorption must be in addition to the complex, narrow intrinsic absorption system seen, would have to almost exactly cancel the intrinsic broad emission (which might mean that it is the first known example of a broad absorption line trough which covers the BLR but not the continuum source), and must be present in \\ovi, \\Nv\\ and \\lya\\ but not in \\SIiv\\ or in \\civ, which is quite unusual. In any case, it seems that some {\\em combination} of unusual parameters is required to explain \\nlya, which helps explain why its spectrum is one of a kind. The X-ray spectrum of \\nlya\\ should consist of a typical power-law plus absorption if the postulated highly-ionized, smooth absorption is real. On the other hand, if its line ratios are largely intrinsic and the trends of \\civ/\\lya\\ with other observables found by \\markcite{bev99}{Wills} {et~al.} (1999) can be extrapolated over an order of magnitude to apply to \\nlya, it should have a very soft X-ray spectrum, a very broad H$\\beta$ line, weak optical \\feii\\ emission, and strong \\oiii\\ emission. \\markcite{bev99}{Wills} {et~al.} (1999) suggest all these trends may derive from a small Eddington parameter (the accretion rate relative to the Eddington rate). While \\nlya\\ may not provide direct insight on the typical quasar BLR if it is indeed a high-density BLR illuminated by an unusual SED, it does help delineate the range of physical parameter space which BLRs occupy and which must therefore be incorporated into BLR models. For example, quasar broad-line \\civ/\\lya\\ ratios may be strongly affected by free-free heating of the BLR, but as \\markcite{fer99}{Ferland} (1999) point out, very little work has been done to investigate the effects of free-free heating on quasar spectra, despite the fact that it can have a greater effect on the observed spectrum than the incident X-ray continuum. The most useful future observation of \\nlya\\ would probably be flux-calibrated spectroscopy at higher resolution to better determine the true profiles of the emission lines, and thus the presence or absence of smooth absorption, by reducing the confusing effects of the many narrow absorption lines present (and also to enable physical modeling of the narrow, associated absorbers). It would also be valuable to obtain an X-ray hardness ratio measurement and near-IR spectroscopy of \\mgii, H$\\gamma$ and H$\\beta$. And if \\nlya\\ has a BLR exposed to a power-law continuum extending from X-ray wavelengths to $\\geq$200\\,\\micron, or $\\geq$800\\,\\micron\\ observed, SCUBA or {\\em Spitzer} photometry at those wavelengths could search for that continuum directly (with the caveat that there is good evidence that the continuum illuminating the BLR in quasars can be different from the continuum seen along our lines of sight to individual quasars; \\markcite{kfb97}{Korista}, {Ferland}, \\& {Baldwin} 1997b)." }, "0405/astro-ph0405525_arXiv.txt": { "abstract": "Using the RHESSI satellite as a Compton polarimeter, a recent study claimed that the prompt emission of GRB021206 was almost fully linearly polarized. This was challenged by a subsequent reanalysis. We present an novel approach, applying our method to the same data. We identify Compton scattering candidates by carefully filtering events in energy, time, and scattering geometry. Our polarization search is based on time dependent scattering rates in perpendicular directions, thus optimally excluding systematic errors. We perform simulations to obtain the instrument's polarimetric sensitivity, and these simulations include photon polarization. For GRB021206, we formally find a linear polarization degree of $\\Pi_{GRB}= (41 ^{+57}_{-44})$\\%, concluding that the data quality is insufficient to constrain the polarization degree in this case. We further applied our analysis to GRB030519B and found again a null result. ", "introduction": "\\label{sec:intro} One of the most outstanding problems in present-day models of the energy release in $\\gamma$-ray bursts (GRB) relates to the conversion of the liberated energy into observed electromagnetic radiation. In relativistic fireball models, the pressure of the photon-lepton plasma itself leads to relativistic expansion (e.g., \\citealt{piran99}). Baryons are then accelerated through their coupling to the electrons. At a more basic level, the origin of the expanding fireball requires a definitive source of energy release that may itself be a source for mass acceleration. \\citet{woosley99} and \\citet{macfadyen99} discuss the formation of relativistically expanding jets from hyperaccreting, stellar-mass black holes that are formed as a consequence of iron core collapse of a rotating, massive ($>30M_{\\odot}$) star. Alternative models for the extraction of rotational (disk and black hole) energy and its conversion to expanding shells or jets involve magnetohydrodynamic processes, including reconnection and dynamo operation in the accretion disks (e.g, \\citealt{blandford77, galeev79, thompson96, katz94, katz97, meszaros97}). The most direct evidence for the presence of magnetic fields in GRB comes from their emission spectrum which is now widely interpreted in terms of synchrotron emission from relativistic electrons. Whether there is a direct connection between magnetic fields produced in the immediate environment of the GRB or the precursor star and the fields that are, at much larger distances, responsible for the observed synchrotron emission is not clear. \\citet{medvedev99} argue that the required field strengths {\\it for GRB afterglow shocks} exceed those to be expected from dragging a pre-existing progenitor field along the expanding shell, and that they also exceed field strengths that could be produced by compression of interstellar magnetic fields. A field-generating mechanism intrinsic to the shocks would thus be required. The situation is, however, less clear for the prompt GRB emission. The recent report \\citep{CB2003} of strong linear polarization in the {\\it prompt} $\\gamma$-ray emission from GRB021206 \\citep{grb021206} therefore stirred some excitement. The observations were performed with the RHESSI satellite \\citep{Lin02}. From extensive modeling, the authors derived a polarization degree of $\\Pi = (80\\pm 20)$\\%, compatible with the assumption of maximum polarization. The emitting electrons are commonly thought to be accelerated in collisionless shocks by the Fermi mechanism to a energy distribution, ${dN/dE} \\propto E^{-p}$, with typically $p \\approx 2$. In that case, standard synchrotron theory predicts that the intrinsic polarization degree $\\Pi$ is a function of $p$, namely $\\Pi = (p+1)/(p + 7/3)$ \\citep{rybicki79}. This value therefore constitutes a maximum for a homogeneous magnetic field. For $p = 2$, one thus finds $\\Pi_{\\rm max} \\approx 0.7$. Tangled magnetic fields with different field vectors relative to the line of sight will in general reduce $\\Pi$. The implications of these observations are far-reaching: Not only does this observation further support the synchrotron radiation model, it also seems to require nearly homogeneous magnetic fields over the source visible to the observer (a solid angle of $\\approx \\Gamma^{-2}$). \\citet{CB2003} argue that magnetic fields dragged by the expanding shell from the surface of the exploding object are too weak to produce the prompt GRB, requiring additional turbulent shock-generated fields that, however, would produce unpolarized emission. Therefore, they suggest that the magnetic fields produced in the central engine are responsible for driving the fireball themselves. Alternatively, post-shock dynamos could generate the highly ordered magnetic fields; but then it will have to be shown that the instabilities occur on large spatial scales \\citep{CB2003}. \\citet{lyut03} calculated the pulse-averaged polarization degree for a relativistically expanding shell that contains a global toroidal field. The synchrotron emission is again assumed to be from power-law distributed electrons. Their calculations predict a maximum of 60\\% for the polarization fraction, depending on the spectral parameters of the GRB. This view is not unequivocal. \\citet{waxman03} points out that a slight off-axis orientation of the relativistically expanding jet produces a strong linear polarization signal even in the presence of random magnetic fields. For further views on this, we refer to the summary in Lyutikov et al. (2003) and references cited therein. In any case, the prospect of diagnosing magnetic field structures in the very emission region of the prompt GRB deserves upmost attention. \\citet{RF2003} have reanalyzed the RHESSI data of GRB021206 and cast serious doubts on the polarization measurements reported earlier by \\citet{CB2003}. \\citet{RF2003} conclude that the signal reported by \\citet{CB2003} is either spurious or is not related to polarization, and they claim that no statistically meaningful statement can be made for the degree of linear polarization for this particular GRB. We have reanalyzed GRB021206 and find problematic issues in both previous analyses. We find, in line with \\citet{RF2003}, clear evidence that much of the signal claimed by \\citet{CB2003} is induced not by source polarization but by accidental coincidences of two unrelated photons arriving at the same time in two different detectors. We essentially confirm the null results of \\citet{RF2003}, but we present a polarization analysis that compares simultaneous scattering rates of orthogonal detector pairs and thus does not need a complicated normalization. Additionally, after implementing photon polarization in the GEANT3 software, we made extensive simulations of RHESSI's response to a fully polarized GRB. We also add one further GRB to this analysis, again finding no statistically significant evidence for non-zero polarization. The structure of the paper is as follows: We start with a description of the RHESSI satellite and its relevant features for polarization analysis in \\S \\ref{sec:rhessi}. Next, we present the RHESSI data of GRB021206 in \\S \\ref{sec:data}. In our data analysis (\\S \\ref{sec:method}), we first select coincidence events (\\S \\ref{sec:c_types}), before we search for a polarization signal in \\S \\ref{sec:polanal}. Then, in \\S \\ref{sec:simul}, we present our simulations of a fully polarized GRB. The results of the polarization analysis of GRB021206 are presented in \\S \\ref{sec:results}, together with the results from our simulation of a fully polarized GRB. After a quick look at GRB030519B - another candidate for polarization analysis - we compare our analysis with previous works (\\S \\ref{sec:comp}) and discuss in general the suitability of RHESSI as a GRB-polarimeter in \\S \\ref{sec:req}. ", "conclusions": "The possibility that prompt GRB emission is strongly linearly polarized at $\\gamma$-ray energies deserves attention with the advent of detectors that can potentially measure such polarization. Polarization carries fundamental information on the orientation of magnetic fields in the source and eventually helps confine the magnetic field generation mechanism. Its detection is difficult due to various interactions in presently available detector systems. The RHESSI satellite can in principle be used as a polarimeter by using the direction dependence of Compton scattering in conjunction with the rotation of the satellite. The effects are subtle, however, and require accurate knowledge of the mass distribution of the satellite and the detector geometry. The methods essentially analyze Compton scattered GRB photons that induce coincidences in detector pairs. A principal difficulty is the separation of such events from accidental and background coincidences in the same detector pairs. Two previous publications \\citep{CB2003, RF2003} have addressed this problem, arriving at essentially opposite conclusions for the same observed GRB. Whereas \\citet{CB2003} claim to detect maximum polarization ($\\Pi = 80\\pm 20$\\%), \\citet{RF2003} challenge this result and find no significant constraint for the polarization degree. We have revisited the problem of GRB polarization measurements with the RHESSI satellite. Our basic test case is GRB021206 used by \\citet{CB2003} and \\citet{RF2003} for their respective analyses. By applying well justified selection criteria for coincidences in energy, time, and scattering angle, we found $N_{C} = 770 \\pm 49$ Compton scattering candidates. By plotting these 770 events as ``coincidence light curves'' for the four different scattering directions, we could define an asymmetry and search for a possible polarization signal. We compare the result of our polarization analysis with simulations of a 100\\% polarized GRB. We cannot reject the null hypothesis that the burst is unpolarized, but neither can we significantly detect any non-zero polarization. The maximum possible synchrotron polarization degree would, for the measured spectral index of the burst, be of order 70\\%, fully compatible with the data. We conclude that {\\em no statement on $\\gamma$-ray polarization can be made for GRB021206}. Our result contradicts the statements of \\citet{CB2003}, who find more than 9000 Compton scattering candidates ($N_C$) and claim to see a polarization signal at the 5.7$\\sigma$ level. The main problem in their analysis is the number of accidental coincidences ($N_{acc}$) that is, in our view, not determined correctly. \\citet{CB2003} obtain their signal by comparing a measured ``angular light curve'' with simulations of an unpolarised GRB. If they do not use a correct $N_{acc}/N_C$-ratio in their simulation, then the presented simulated data points cannot be trusted. Our result agrees in many aspects with the reanalysis presented by \\citet{RF2003}. By making a more sophisticated coincidence selection, we obtain much smaller errors, however. A few points of caution in the polarization analysis of \\citet{RF2003} have been mentioned in our presentation. We also analyzed the strong GRB030519B but found fewer Compton scattering candidates than for GRB021206 and, again, no statement about the polarization degree of GRB030519B can be made. We therefore conclude that RHESSI has not yet measured the degree of polarization of any observed GRB. Implications on magnetic field orientation based on RHESSI results are therefore premature, and the physics of magnetic field generation and structure in the $\\gamma$-ray source of GRB must rely on alternative information for the time being. On the other hand, our analysis suggests that at least in principle, RHESSI does offer the capability of measuring polarization. A GRB with comparable fluence as GRB021206, but lasting $\\approx 10\\,$s, would be required to significantly distinguish maximum polarization of a GRB from zero polarization. In addition, a GRB suitable for polarization analysis should come from a direction close to RHESSI's rotation axis." }, "0405/astro-ph0405477_arXiv.txt": { "abstract": "The HIDEEP survey (Minchin et al.~2003) was done in an attempt to find objects having low inferred neutral hydrogen column densities, yet they found a distribution which was strongly peaked at $10^{20.65}$ cm$^{-2}$. In an attempt to understand this distribution and similar survey results, we model HI profiles of gas discs and use simple simulations of objects having a wide range of HI properties in the presence of an ionizing background. We find that inferred column density ($N_{HI}^o$) values, which are found by averaging total HI masses over some disc area, do not vary strongly with central column density ($N_{max}$) for detectable objects, so that even a population having a wide range of $N_{max}$ values will give rise to a strongly peaked distribution of $N_{HI}^o$ values. We find that populations of objects, having a wide range of model parameters, give rise to inferred column density distributions around $10^{20.6\\pm 0.3}$ cm$^{-2}$. However, populations of fairly massive objects having a wide range of central column densities work best in reproducing the HIDEEP data, and these populations are also consistent with observed Lyman limit absorber counts. It may be necessary to look two orders of magnitude fainter than HIDEEP limits to detect ionized objects having central column densities $<10^ {20}$ cm$^{-2}$, but the inferred column densities of already detected objects might be lower if their radii could be estimated more accurately. ", "introduction": "Understanding the properties of dwarf galaxies, large diffuse galaxies, and any clouds of similar mass is important in understanding the formation of galaxies. For example, Cold Dark Matter theory predicts the existence of a population of low-mass satellite galaxies (e.g. Moore et al. 1999; Klypin et al. 1999). Furthermore, studying the properties of such objects is important in understanding the nature of Ly$\\alpha$ absorbers and metal line absorbers such as weak MgII systems (Rigby, Charlton, \\& Churchill 2002). Any such objects which have yet been undetected may have a different range of averaged neutral hydrogen column densities from that of the known population of galaxies. Gas having a wide range of neutral column density ($N_{HI}$) values has been observed as Ly$\\alpha$ absorption at low redshifts (for example, Bahcall et al.~1996) where absorption lines shortward of Ly$\\alpha$ emission in quasar spectra arise from lines of sight through intervening gas between us and the quasar, and $N_{HI}$ ranges from $<10^{12}$ cm$^{-2}$ to $\\sim 10^{21}$ cm$^{-2}$. Larger amounts of gas with $N_{HI}\\gtorder 10^{19}$ cm$^{-2}$ can also be observed more directly as 21 cm emission in the local universe. The strongest `damped' Ly$\\alpha$ absorbers, with $N_{HI} > 10^{20.3}$ cm$^{-2}$, are often found to arise in lines of sight through galaxies including several low surface brightness (LSB) and dwarf galaxies (Cohen 2001; Turnshek et al.~2000; Bowen, Tripp, \\& Jenkins 2001). Yet the somewhat weaker Lyman limit systems ($N_{HI} > 10^{17.2}$ cm$^{-2}$), the column densities of which are more difficult to measure accurately, have long been thought to arise in lines of sight through luminous galaxies (Bergeron \\& Boiss\\'e 1991; Steidel 1995). Some weaker Ly$\\alpha$ forest absorbers are thought to arise in small amounts of intergalactic gas (Dav\\'e et al.~1999), while some could arise in gas surrounding galaxies (e.g. Chen et al.~2001; Linder 1998; 2000). A recent HI survey (HIDEEP; Minchin 2001; Minchin et al. 2003) was capable of detecting objects with inferred neutral hydrogen column densities ($N_{HI}^o$) as low as $4\\times 10^{18}$ cm$^{-2}$ for galaxies having velocity width $\\triangle V=200$ km s$^{-1}$, assuming that a galaxy with suitable properties fills the telescope beam. Yet they failed to find anything with $N_{HI}^o<10^{20}$ cm$^{-2}$. Other HI surveys have also found that galaxies show little variation in column densities averaged over some radius (Zwaan et al.~1997), although the integration times may not be long enough to detect low column density galaxies in such surveys, as discussed by Minchin et al.~(2003). These HI surveys are limited by flux, rather than column density, when detecting faint objects, and the column density of the detected objects is uncertain given that the sources are generally unresolved. However there is a limit on column density in a survey such as HIDEEP in the sense that a resolved, low column density object could fill the beam, although such objects are not often seen. Rosenberg \\& Schneider (2003) found that their sample of HI-selected galaxies obey a relationship between HI cross section and HI mass, which is equivalent to having fairly constant averaged column densities. They plot, in their first figure, the disc areas $A_{DLA}$, where $N_{HI}>2 \\times 10^{20}$ cm$^{-2}$ and thus where damped Ly$\\alpha$ absorbers can arise, versus the HI mass ($M_{HI}$) for a sample of HI selected galaxies. Some scatter is seen in the log-log plot, yet they can easily fit a line having a slope of about one. Thus they find $\\log(A_{DLA})=\\log(M_{HI})-6.82$, which would imply that galaxies having a wide range of mass and HI sizes all have area-averaged column densities of around $8\\times 10^{20}$ cm$^{-2}$, where the displayed points are all within about $0.8$ orders of magnitude from the fitted line. Similar correlations between HI size and HI mass have also been seen by Giovanelli \\& Haynes (1983) and Verheijen \\& Sancisi (2001), and a correlation between HI mass and optical sizes of galaxies has also been seen by Haynes \\& Giovanelli (1984). Other surveys, capable of detecting low HI mass objects at various sensitivities, including some directed toward detecting extragalactic high velocity clouds (HVCs) (Blitz et al.~1999; Charlton, Churchill \\& Rigby 2000; Davies et al.~2002) have been largely unsuccessful at finding objects with low HI masses (de Blok et al.~2002; Zwaan \\& Briggs 2000; Dahlem et al. 2001; Zwaan 2001; Verheijen et al.~2000). On the other hand, some very faint optical sources have been found to be rich in gas (Davies et al.~2001), and there is theoretically no reason to expect every HI cloud to be capable of forming large amounts of stars. Furthermore, small HVCs with peak $N_{HI}\\sim 6\\times 10^{18}$ cm$^{-2}$ are being detected around our Galaxy (Hoffman, Salpeter, \\& Pocceschi 2002) and around M31 (Thilker et al.~2004). One suggested explanation for the lack of low column density detections in the HIDEEP survey is that the gas is hidden in 'frozen discs' (Minchin et al.~2003) where the 21 cm transition is not excited to a spin temperature above the cosmic background (Watson \\& Deguchi 1984). A second possible explanation for the lack of low column density detections is that the gaseous discs become highly ionized at a disc radius not far beyond that where $N_{HI}=10^{20}$ cm$^{-2}$, so that the average inferred value remains above $10^{20}$ cm$^{-2}$. The ionization of outer galaxy discs by a background of Lyman continuum photons was suggested and modelled first by Bochkarev \\& Sunyaev (1977) and later by Maloney (1993), Dove \\& Shull (1994a), and Corbelli \\& Salpeter (1993) in order to explain the sudden truncations seen in carefully observed spiral galaxy discs. Since then ionized gas has been detected in H$\\alpha$ emission using a Fabry-Perot `staring technique' (Bland-Hawthorn et al.~1994) beyond the HI edges of several nearby galaxies (Bland-Hawthorn, Freeman \\& Quinn 1997; Bland-Hawthorn 1998). The ionizing background has been measured most recently at low redshifts by Scott et al. (2002), who find $J(912$\\AA$)=7.6^{+9.4}_{-3.0}\\times10^{ -23}$ erg cm$^{-2}$ s$^{-1}$ Hz$^{-1}$ sr$^{-1}$. Gas in the ionized parts of outer galaxy discs is likely to give rise to at least some Ly$\\alpha$ absorption (Linder 1998; 2000), and some variations will arise in the column density value at which HI discs are truncated as a result of fluctuations in the ionizing background radiation (Linder et al.~2003). Ionized gas clouds cannot correctly be referred to as undetected 'HI clouds' (although current HI structures may have been ionized in the recent cosmological past). However structures containing ionized gas are interesting and relevant to the galaxy formation process. For example, ionized gas contains enough neutral atoms to give rise to all of the Lyman alpha absorbers (except for the damped ones), and is thus, in principle, detectable in deep HI observations. HVCs may also contain mostly ionized gas. It is unknown whether massive clouds exist far from luminous galaxies, although not all of the absorbers arise close to galaxies (Stocke et al.~1995). Ionized gas clouds may have small regions containing HI clouds if the gas is sufficiently clumpy. In this paper, we wish to understand the observed lower limits in averaged column densities, and to constrain the properties of any objects that could be going undetected in HI surveys as a result of photoionization. Section 2 discusses the modelling of HI discs and calculation of column densities from HI observations. Section 3 describes the method used to model galaxy and cloud HI profiles and simulate populations of objects having a wide range of properties in the presence of an ionizing background. The results of such simulations are discussed in Section 4. The value for the Hubble constant is assumed to be $H_0=80$ km s$^{-1}$ Mpc$^{-1}$. ", "conclusions": "Most galaxies have inferred column densities around $10^{20.6\\pm 0.3}$ cm$^{-2}$ because inferred column densities are found by averaging column density profiles, which are exponential or similar, over a radius where the minimum column density is $\\sim 10^{20}$ cm$^{-2}$. Ionization plays some role in making lower column density objects undetectable, including those without substantial optical counterparts. However inferred column density distributions tell us little about the distribution of central column densities in galaxies and clouds. Ionization by the background of ultraviolet photons will strongly affect the amount of neutral gas remaining, and thus the HI flux detected, in objects having low hydrogen column densities, if such objects having sizes comparable to galaxies exist. Typical HI fluxes are reduced, as a result of ionization, by a factor of $\\sim 100$ for galaxies having peak column densities $N_{max}\\sim 10^{19.5}$ cm$^{-2}$ compared to those with $N_{max}\\sim 10^{20}$ cm$^{-2}$, even if the lower column density galaxies are extended in size and just as massive as the higher column density galaxies. We do not always know the central column densities of the faintest HI sources, but the detected inferred column densities are also likely to be above $\\sim 10^{20}$ cm$^{-2}$ for most observable galaxies. Inferred column densities are rather weakly related to central column densities for objects having exponential profiles. Furthermore, since HI profiles tend to be mapped out to limiting column densities $\\sim 10^{20}$ cm$^{-2}$, it may be difficult to estimate the radii, and thus the inferred column densities in a consistent manner for objects having lower $N_{max}$ values. For example, if the radii are underestimated, which might be more likely to happen for an extended, diffuse galaxy, the inferred column density could be overestimated. Other selection effects, such as those against objects having low velocity widths, may also be important in understanding the observed distribution of inferred HI column densities. The observed distribution of inferred HI column densities, as seen by Minchin et al.~(2003), can easily be simulated assuming possible populations of galaxies having a wide range of size and central column density distributions, and the simulated distributions are similar to the HIDEEP distribution for a wide range of model parameters. (Thus the 'Frozen Disc' hypothesis of Minchin et al.~2003 seems to be unnecessary in explaining these observations.) However, we are thus given little constraint on the properties of gas rich objects which have so far escaped detection in the deepest HI surveys. Given the effects of ionization, we are unable to rule out the existence of undetected populations of very faint dwarf galaxies or giant gas clouds, as long as they have low central column densities. Such objects could make some contribution to Ly$\\alpha$ absorption, although a more reasonable number of Lyman limit systems arises if galaxies have a wide, rather than narrow, range of central column densities. The ionizing background radiation is more intense at redshifts around 1 or 2 than at redshift zero (Haardt \\& Madau 1996), and therefore some of the apparently younger galaxies, such as LSB galaxies, may have been ionized at these redshifts if they have lower central column densities (de Blok et al.~1996), thus slowing their evolution. Ionization may have also affected the formation of dwarf galaxies in certain environments at high redshifts (Efstathiou 1992; Tully et al.~2002), as less dense environments are more likely to be optically thin to ionizing radiation when the dwarf galaxies formed. Thus dwarf galaxies may have formed more easily in rich clusters such as Virgo (Sabatini et al.~2003) and Fornax (Kambas et al.~2000) than in more diffuse clusters such as Ursa Major (Trentham \\& Tully 2002) and other environments (Roberts et al.~2003). Understanding the role that ionization plays is thus important in testing Cold Dark Matter scenarios and other theories related to galaxy formation." }, "0405/astro-ph0405194_arXiv.txt": { "abstract": "PSR~J0737$-$3039A is a millisecond pulsar with a spin period of 22.7 ms included in a double-neutron star system with an orbital period of 2.4 hrs. Its companion has also been detected as a radio pulsar, making this binary the first known double-pulsar system. Its discovery has important implications for relativistic gravity tests, gravitational wave detection and plasma physics. Here we will shortly describe the discovery of the first pulsar in this unique system and present the first results obtained by follow-up studies. ", "introduction": "Since the discovery of the first binary pulsar (Hulse \\& Taylor 1975), the detection of two active pulsars in the same binary system has been a primary aim of any pulsar survey. We here summarize the basic steps which eventually led to the discovery of this long-sought system and report on the first implications for gravitational waves detection. Papers by Kramer et al. and Manchester et al. (in these proceedings) will give more details on the second discovered pulsar, dealing with the opportunity to use this binary as a magnificent laboratory of relativistic gravity and for investigating magnetospheric processes. ", "conclusions": "" }, "0405/astro-ph0405127_arXiv.txt": { "abstract": "We present an analysis of observations of the bright star Altair ($\\alpha$ Aql) obtained using the star camera on the Wide-Field Infrared Explorer (WIRE) satellite. Although Altair lies within the $\\delta$ Scuti instability strip, previous observations have not revealed the presence of oscillations. However, the WIRE observations show Altair to be a low-amplitude ($\\Delta m < 1$~ppt) $\\delta$ Scuti star with at least 7 modes present. ", "introduction": "Delta Scuti stars are a variety of pulsating variable star located within the classical instability strip of the HR diagram (Rodriguez \\& Breger 2001). They inhabit the region between $3.8 < \\log \\rm T_{eff} < 3.95$ and $0.6 < \\log \\rm (L/L_\\sun) < 2.0$, and display periods ranging from 0.02 to 0.3 days. Some belong to the high-amplitude $\\delta$ Scuti (HADS) class, which display $V$ amplitudes in excess of 0.3 mag, and generally oscillate in radial modes, while the lower-amplitude members of the class have more complex frequency structure, typically showing numerous nonradial modes. Altair ($\\alpha$ Aql) is an A7 IV-V main sequence star (Johnson \\& Morgan 1953) and is the 12th brightest star in the sky (V = 0.755, Cousins 1984). While it lies inside the instability strip, we are unaware of any reports of photometric variability. Erspamer \\& North (2003) report $\\rm T_{eff} = 7550 K$ and $\\log g = 4.13$, and Zakhozhaj (1979, see also Zakhozhaj \\& Shaparenko 1996) derives a mass of $1.75 M_\\sun$ and a radius of $1.58 R_\\sun$ based on photometry. The Hipparcos parallax for Altair is $194.44 \\pm 0.94 \\rm~ mas$, giving a distance of $5.143 \\pm 0.025 \\rm~pc$. Combined with the observed V magnitude and the bolometric correction from Flower (1996) gives $\\rm M_{bol} = 2.18$. Altair is a rapid rotator, with spectroscopically derived $v \\sin i$ values in the literature ranging from $190 \\rm~km~s^{-1}$ to $250 \\rm~km~s^{-1}$ (Royer et al. 2002). Altair is nearby enough to make direct measurement of its diameter possible. Richichi et al. (2002) reported a diameter of $3.12 \\rm ~mas$, corresponding to a radius of $1.72 \\rm R_\\sun$. Recent interferometric observations (Van Belle et al. 2001) have established that Altair is oblate, with equatorial diameter of $3.46 \\rm ~mas$ (corresponding to a radius of $1.9 \\rm R_\\sun$) and polar diameter of $3.037 \\rm ~mas$, for an axial ratio $a/b = 1.14 \\pm 0.029$. Van Belle et al. also derive a value for $v \\sin i$ of $210 \\pm 13 \\rm~ km~s^{-1}$. Altair's absolute V magnitude is +2.22 (based on the Hipparcos distance) and its $\\rm (b - y)_0 = 0.137$ (Hauck \\& Mermilliod 1998; Nekkel et al. 1980), which places it well inside the $\\delta$ Scuti instability strip (Breger 1990). However, like many such objects, Altair has never been detected to oscillate. It is as yet unclear whether these non-oscillating stars within the instability strip are truly non-oscillating, or whether the non-detections are instead due to either periods or amplitudes which are undetectable from the ground; some indications are that slow rotation is a necessary condition for large oscillation amplitudes (Breger 1982, Solano \\& Fernley 1997, but for an alternate view see Rasmussen et al. 2002). ", "conclusions": "We have determined that the bright A7~IV-V star Altair is an oscillating variable of the $\\delta$ Scuti type, and that its three most significant oscillation frequencies correspond well to the fundamental radial mode and its first two overtones. Four additional frequencies are present at lower amplitudes, and these may represent nonradial modes. Our result suggests that oscillation may in fact be present in many putatively ``non-oscillating'' stars in the instability strip. On the basis of the observed frequencies and the model fits, several specific points are clear: \\begin{enumerate} \\item As suggested by the theoretical P-L relation of Breger \\& Bregman (1975), the highest-amplitude mode is the radial fundamental mode. The amplitude of this oscillation, 0.5~ppt, is several times smaller than is typically detected in non-HADS $\\delta$ Scuti stars from ground-based observations. The first and second overtones are also apparent. \\item It has been suggested that (Breger 1982) that slow rotation is a necessary (but not sufficient) condition for the development of large-amplitude oscillations in $\\delta$ Scuti stars. Our result supports this contention. \\item A number of authors (see, {\\it e.g.} Rasmussen et al. 2002 and references therein) have tried to understand the excitation mechanism in $\\delta$~Scuti stars by searching for systematic differences between variable and non-variable stars. Our result suggests that such searches may be complicated by the fact that ``non-variable'' stars may in fact be variable at levels undetectable from the ground. In fact, the division between ``variable'' and ``non-variable'' stars in the instability strip may truly be an observational (detection technology-limited) one, with the objects in question occupying a broad continuum of photometric amplitudes. \\item Non-rotating models should not be expected to yield accurate representations of oscillation frequencies or ratios in rapidly rotating $\\delta$ Scuti stars. There is a significant need for calculations of oscillation frequencies that incorporate the effects of rotation. \\end{enumerate}" }, "0405/astro-ph0405582_arXiv.txt": { "abstract": "We present Very Small Array (VSA) observations (centred on $\\approx$ 34\\,GHz) on scales $\\approx$ 20 arcmin towards a complete, X-ray--flux--limited sample of seven clusters at redshift $z<0.1$. Four of the clusters have significant Sunyaev-Zel'dovich (SZ) detections in the presence of CMB primordial anisotropy. For all seven, we use a Bayesian Markov-Chain-Monte-Carlo (MCMC) method for inference from the VSA data, with X-ray priors on cluster positions and temperatures, and radio priors on sources. In this context, the CMB primordial fluctuations are an additional source of Gaussian noise, and are included in the model as a non--diagonal covariance matrix derived from the known angular power spectrum. In addition, we make assumptions of $\\beta$--model gas distributions and of hydrostatic equilibrium, to evaluate probability densities for the gas mass ($M_{\\rm{gas}}$) and total mass ($M_{\\rm{r}}$) out to $r_{200}$, the radius at which the average density enclosed is 200 times the critical density at the redshift of the cluster. This is further than has been done before and close to the classical value for a collapsed cluster. Our combined estimate of the gas fraction $(f_{\\rm{gas}}=M_{\\rm{gas}}/M_{\\rm{r}})$ is $0.08^{+0.06}_{-0.04}h^{-1}$. The random errors are poor (note however that the errors are higher than would have been obtained with the usual chi-squared method on the same data) but the control of bias is good. We have described the MCMC analysis method specifically in terms of SZ but hope the description will be of more general use. We find that the effects of primordial CMB contamination tend to be similar in the estimates of both $M_{\\rm{gas}}$ and $M_{\\rm{r}}$ over the narrow range of angular scales we are dealing with, so that there is little effect of primordials on $f_{\\rm{gas}}$ determination. Using our $M_{\\rm{r}}$ estimates we find a normalisation of the mass -- temperature relation based on the profiles from the VSA cluster pressure maps that is in good agreement with recent $M-T$ determinations from X-ray cluster measurements. ", "introduction": "\\label{sec:introduction} Galaxy clusters have long been thought to provide a faithful sample of cosmic baryonic matter (see e.g. \\cite{COS/Whi++93}, \\cite{Evrard}). One quantity often calculated and assessed in such work is the gas fraction $f_{\\rm{gas}}$, which is defined as the (baryonic) gas mass over the total (baryonic plus dark matter) mass of the cluster. We here present Sunyaev Zel'dovich (SZ) (\\cite{SZ}, see also e.g. \\cite{Birkinshaw1999}, \\cite{Carlstrom_review}) observations of a sample of clusters, from which we infer $f_{\\rm{gas}}$. Our random errors are high but the sample is complete, the redshifts deliberately low, and we are able to estimate $f_{\\rm{gas}}$ out to radii at which the overdensity of the enclosed region is close to the classical value of 178 for a collapsed object (see e.g. \\cite{Peacock}). First we review some of the existing $f_{\\rm{gas}}$ measurements. A popular route in investigating cosmic baryonic matter is the detailed study of the X-ray emission from cluster gas. For example, in an investigation based on \\emph{ROSAT} PSPC data (\\cite{fabian_fgas}), a sample of 36 clusters of redshift $0.05 \\le z \\le 0.44$ was used to measure $f_{\\rm{gas}}$. Assumptions of isothermality and hydrostatic equilibrium were required. The resulting $f_{\\rm{gas}}$ distribution (within $r_{500}$, that is, where the mean density inside this radius is 500 times the critical density at the redshifts of the clusters) was centred on a value $f_{\\rm{gas}}(r_{500})=0.168h_{50}^{-1.5}$. Values for individual clusters were found to vary between 0.101 and 0.245. \\cite{Mohr} also analysed PSPC data on 45 X-ray selected clusters, finding a mean $f_{\\rm{gas}}(r_{500})$ of $0.212h_{50}^{-1.5}$ in a subsample of 27 clusters hotter than 5 keV. \\cite{IOA1}, following a similar route (supplemented by gravitational lensing information on the total mass) with \\emph{Chandra} imaging spectrometer data find, for a set of six clusters with $0.103 \\le z \\le 0.461$, a mean $f_{\\rm{gas}}$ within $r_{2500}$ of $0.113 \\pm 0.005h_{70}^{-1.5}$ for a $\\Lambda$--CDM model, a very precise determination with very similar values for each cluster. \\cite{IOA2}, with additional data, investigated the observed change of $f_{\\mathrm{gas}}$ with cosmology. Studies making use of the SZ effect have potential advantages for gas and gravitational potential measurements (where the potential is obtained via calculation of the total mass). The X-ray signal is proportional to $n_{\\rm{e}}^2$ (where $n_{\\rm{e}}$ is electron density), while the SZ signal is proportional to $n_{\\rm{e}}$. This means that SZ is less biased to concentration and can constrain clumping. Although X-ray telescopes achieve excellent signal to noise, they are restricted to observing the denser, inner regions of a cluster (e.g out to $r_{2500}$). With SZ it is possible to measure $n_{\\rm{e}}(r)$ over a larger range of r (e.g. close to the virial radius) as less dynamic range is required. \\cite{Myers_obs} used the OVRO 5.5m telescope to observe the SZ effect in 3 clusters at 32\\,GHz. With the addition of the Coma cluster (observed by \\cite{Herbig}), they obtain a gas fraction of $f_{\\rm{gas}}=0.061\\pm{0.011}h^{-1}_{100}$ This sample of objects lies in the redshift range $0.023 \\leq z \\leq 0.0899$, and includes three clusters which we also present here. (\\cite{Mason_obs} extend the sample to seven clusters, incorporating a further two discussed in this paper. The data were used to calculate $H_0$.) \\cite{grego} used the OVRO and BIMA arrays to make SZ observations of galaxy clusters at 30\\,GHz. The data were used to infer the gas mass and total mass, thus constraining $f_{\\rm{g}}$ (within $r_{500}$) in 18 X-ray selected clusters in the redshift range $0.171\\leq z \\leq 0.826$. The mean value obtained for the full sample was $f_{\\rm{gas}}=0.081^{+0.009}_{-0.011}h^{-1}_{100}$. In addition, a `fair' subsample is defined as the five most X-ray luminous clusters in the EMSS sample. These objects have redshift $0.328\\leq z \\leq 0.826$, and together give a mean gas fraction $f_{\\rm{gas}}=0.089^{+0.018}_{-0.019}h^{-1}_{100}$. One of the aims of the VSA project (\\cite{paper1}, \\cite{paper2}, \\cite{paper3}, \\cite{paper4}, \\cite{paper5}, \\cite{anze} \\cite{clive}, \\cite{rafa}) has been to image nearby, massive clusters in SZ. The VSA baselines at $\\approx34\\rm{\\,GHz}$ couple well to the angular scales of such clusters. Here we describe SZ observations and cluster--parameter inferences of an X-ray selected, complete sample of seven clusters, with redshift $0.023 \\leq z \\leq 0.098$ and median 0.075. The age of the Universe at z = 0.075 is 1.7 times its age at $z = 0.55$. The importance of low--z work is illustrated by the following two points: \\begin{itemize} \\item {The low redshifts of the clusters mean that they have particularly good X-ray data, and one can be reasonably confident that bright X-ray selected complete samples are in fact complete.} \\item {Since clusters grow under gravity, then on average low redshift clusters should be more evolved than those at higher redshift. Comparison of, for example, $f_{\\rm{gas}}$ in low- and high-$z$ samples is important. (Of course, we do not know how big the samples have to be to encompass meaningful averages).} \\end{itemize} One immediate difficulty on these angular scales is contamination by CMB primordial anisotropy. At the start of this VSA observational programme, it was evident that we needed an analysis method that would apply the inference process correctly and would properly cope with error distributions in low signal--to--noise situations. There is the additional difficulty of dealing with (potentially variable) radio sources at 34\\,GHz. This could be especially problematic where sources are in the clusters themselves rather than in the background: the low redshifts of the clusters imply such sources may be very bright. Accounting for these effects correctly necessitates the exploration of the posterior probabilities of the parameters of a $\\beta$--model for the gas distribution given the VSA visibilities, receiver noise, the CMB and radio sources. The method must also incorporate prior knowledge on e.g. the cluster positions from X-rays, and on source fluxes in a way which can cope with variability. We assume isothermality, and that the clusters are well described by hydrostatic equilibrium. We use a Markov Chain Monte Carlo (MCMC) sampler (\\texttt{BayeSys}) for an acceptable combination of speed and accuracy. In section 2 we briefly describe the relevant features of the VSA. In section 3 we present the sample, outline the data reduction pipeline and describe our strategy for dealing with radio sources. In section 4 we present our results, and attempt to describe the Bayesian analysis method in non-specialist terms. We make concluding comments in section 5. ", "conclusions": "We have investigated with the VSA Extended Array at $\\approx34$\\,GHz the SZ effects towards seven nearby clusters that form a complete, X-ray--flux--limited sample. \\begin{enumerate} \\item{Four of the clusters (Coma, A1795, A478, A2142) show SZ effects in the map plane on scales of $\\approx$20 arcmin of typically 6$\\sigma$.} \\item{There is significant detection of CMB primordial structure at this resolution, which is the likely cause of the three non-detections (A399, A401, A2244). We have analysed the data in the $uv$--plane, with X-ray priors on positions and gas temperatures and radio priors on the sources, using MCMC to estimate key cluster parameters in the context of a $\\beta$--model for the gas distribution. In this context, the CMB primordial fluctuations are an additional source of Gaussian noise, and are included in the model as a non--diagonal covariance matrix derived from the known angular power spectrum. We use the SZ data (plus the priors) to give both the gas mass and, under the assumption of hydrostatic equilibrium, the total mass. Although the data have high random errors, the use of Bayesian methods, probability density functions and marginalisation prevents bias in the results.} \\item{The degeneracy is evident between $\\beta$ and core radius as expected for such observations sensitive to SZ over a narrow $\\ell$--range. There are significant measurements of gas fractions in the detected clusters.} \\item{We present a normalisation of the M-T relation derived from our data which we find to be in good agreement with recent X-ray cluster measurements.} \\item{Using the gas fraction probability density function for each cluster, we have produced combined gas fractions for the four detections, for the three non-detections, and for all seven. The Bayesian evidence shows that the first is the correct one to use in the context of trying to measure a low-$z$ global gas fraction. For this we here find $f_{\\rm{gas}} = 0.08^{+0.06}_{-0.04} h_{100}^{-1}$.} \\item{Gas fraction measurement by this SZ--based method is relatively immune from the effect of primordial CMB anisotropy. This is true since the effect on gas mass tends to cancel the effect on total mass on the narrow range of angular scale employed. Simulations show the cancellation to be good for contaminant fluxes of $\\pm50$\\,mJy.} \\end{enumerate} That the analysis method works as well as it does points the way towards analysis of data from upcoming SZ telescopes." }, "0405/astro-ph0405257_arXiv.txt": { "abstract": "{ We present a new analysis on the issue of the location of the observed microlensing events in direction of the Large Magellanic Cloud (LMC). This is carried out starting from a recently drawn coherent picture of the geometrical structure and dynamics of the LMC disk and by considering different configurations for the LMC bar. In this framework it clearly emerges that the spatial distribution of the events observed so far shows a near--far asymmetry. This turns out to be compatible with the optical depth calculated for the LMC halo objects. In this perspective, our main conclusion, supported by a statistical analysis on the outcome of an evaluation of the microlensing rate, is that self lensing can not account for all the observed events. Finally we propose a general inequality to calculate quickly an upper limit to the optical depth along a line of view through the LMC center. ", "introduction": "The microlensing surveys towards the Large Magellanic Cloud (LMC) \\citep{alcock00a,lasserre00} have demonstrated the existence of compact objects that act as gravitational lenses somewhere between us and the LMC. In some cases the distance and the mass of the lenses have been determined, thanks to the proper motion of the lens observed by the Hubble Space Telescope (HST) \\citep{alcock01a,gould04,drake04} or to the additional information carried by binary systems \\citep{alcock00b,alcock01b}. However, these are special cases since for most events only the duration and the position on the sky plane have been measured. These information are not enough to establish definitively if the detected events are really caused by white dwarfs or MACHOs in the halo of the Milky Way (MW), or are due to stars or MACHOs within the LMC itself. The survey of the MACHO team indicates a most probable Galactic halo fraction of 20\\%, with limits of 5\\% to 50\\% at the 95\\% confidence level, assuming that all the events are due to halo lenses. The preferred value for the lens mass is $\\sim$ 0.4 M$_{\\sun}$. This is consistent with the EROS survey results, that are given however as an upper limit for the Galactic halo fraction. An interesting alternative is that of ``self lensing'', where both source and lens belong to the luminous part of the LMC as suggested by \\citet{sahu94} and \\citet{wu94}. However, the initial estimates of the optical depth and microlensing rate were lower than the measured one \\citep{gould95,alcock97a,alcock00a}. The self--lensing explanation has been further analyzed going beyond the hypothesis of a ``simple'' geometry for the LMC with disk and bar coplanar, so that their relative distance would enhance the optical depth and, therefore, the rate. In the model of \\citet{zhao00} the disk and bar stars are on two distinct planes with different inclinations, so that stars on the front plane could lens those in the plane $\\approx$ 1 kpc behind. Besides the morphology, another aspect considered has been the dynamics of the luminous components within the LMC. \\citet{aubourg99}, by using a model which takes into account the correlation between the mass of the lenses in the LMC and their velocity dispersion, have been able to reproduce a self--lensing optical depth, event rate and event duration distribution compatible with the observed ones. Yet, objections to this model have been raised by different authors \\citep{gyuk00,alves00}, especially with respect to the adopted distribution and velocity dispersion of the lensing stars, which seem to be inconsistent with the observations. The analysis of Jetzer et al., (2002, hereafter Paper I) has shown that probably the observed events are distributed among different components (disk, spheroid and galactic halo, the LMC halo and self--lensing). This means that the lenses do not belong all to the same population and their astrophysical features can differ deeply from one another. In this paper we address once more the question of the presence of a self--lensing component within the LMC itself. To this end a correct knowledge of the structure and dynamics of the luminous components (disk and bar) of the LMC is essential. Here we take advantage of some recent studies of the LMC disk (see Sect. \\ref{sec:disk-morphology}), while we allow for different geometries for the still poorly known bar component, to calculate the main microlensing quantities. Moreover, with respect to Paper I, based on the moment method \\citep{derujula}, we perform instead a statistical analysis starting from the differential rate of the microlensing events. The paper is organized as follows: in Sects. \\ref{sec:morphology} and \\ref{sec:modelli} we discuss the LMC morphology and present the models we use to describe the spatial density of the MW halo and of the various components of the LMC. Sect. \\ref{sec:mlpar} is devoted to the calculation of the microlensing quantities, the optical depth and the microlensing rate, as well as to a statistical analysis of the self--lensing events. In Sect. \\ref{sec:asimmetria} we discuss the spatial asymmetry with respect to the line of nodes of the observed microlensing events. An improved inequality for the optical depth for self lensing by a stellar disk is derived in Sect. \\ref{sec:gould}. We conclude in Sect. \\ref{sec:conclusioni} with a summary of our results. ", "conclusions": "\\label{sec:conclusioni} The great interest in the \\emph{location} of the observed microlensing events towards the LMC is motivated by the need to give an answer to the question of their \\emph{nature}. Namely, whether (or not) all the events can be attributed to known (luminous) populations, so to exclude (or not) the possibility for a dark component in the halo in the form of MACHOs. In this paper we are mainly concerned with the possible self--lensing origin of the observed microlensing events. In particular we have considered the results of the MACHO survey. We use the recent picture of the LMC disk given by \\citet{marel02}, and we explore different geometries for the bar component, as well as a reasonable range for the velocity dispersion for the bar population. One interesting feature, essentially linked to the assumed disk geometry, is an evident near--far asymmetry of the optical depth for lenses located in the LMC Halo (this is not expected, with the possible exception of the inner region, for the self--lensing population). Indeed, similarly to the case of M31 \\citep{crotts,jetzer}, and as first pointed out by \\citet{gould93}, since the LMC disk is inclined, the optical depth is higher along lines of sight passing through larger portions of the LMC halo. We show that such a spatial asymmetry, beyond the one expected from the observational strategy alone, is indeed present in the observed events. With the care suggested by the small number of detected events on which this analysis is based, this can be looked at, as yet observed by \\citet{gould93}, as a signature of the presence of an extended halo around the LMC. In the central region the microlensing signatures are strongly dependent on the assumed bar geometry. In particular, we have studied the variation (that can be as large as 50\\%) in the self--lensing optical depth due to the different geometry of the bar. However, the available data do not allow us to meaningfully explore in more detail this aspect. As a further step in our analysis, we have studied the microlensing rate. Keeping in mind \\citet{evans00} observation that the timescale distribution of the events and their spatial variation across the LMC disk offers possibilities of identifying the dominant lens population, we have carefully characterized the ensemble of observed events under the hypothesis that all of them do belong to the self--lensing population. Through this analysis we have been able to identify a large subset of events that can not be accounted as part of this population. The introduction of a non coplanar bar component with respect to the disk turns out to enhance this result. Again, the small amount of events at disposal does not yet allow us to draw sharp conclusions, although, the various arguments mentioned above are all consistent among them and converge quite clearly in the direction of excluding self lensing as being the major cause for the events. Once more observations will be available, as will hopefully be the case with the SuperMacho experiment under way \\citep{stubbs}, the use of the above outlined methods can bring to a definitive answer to the problem of the location of the MACHOs and thus also to their nature." }, "0405/astro-ph0405510_arXiv.txt": { "abstract": "s{ We investigate nucleosynthesis inside both a gamma-ray burst accretion disk and a wind launched from an inner region of the disk using one-dimensional models of the disk and wind and a nuclear reaction network. Far from a central black hole, the composition of accreting gas is taken to be that of an O-rich layer of a massive star before core collapse. We find that the disk consists of five layers characterized by dominant elements: \\nuc{O}{16}, \\nuc{Si}{28}, \\nuc{Fe}{54} (and \\nuc{Ni}{56}), \\nuc{He}{4}, and nucleons, and the individual layers shift inward with keeping the overall profiles of compositions as the accretion rate decreases. \\nuc{Ni}{56} are abundantly ejected through the wind from the inner region of the disk with the electron fraction $\\simeq 0.5$. In addition to iron group, elements heavier than Cu, in particular \\nuc{Cu}{63} and \\nuc{Zn}{64}, are massively produced through the wind. Various neutron-rich nuclei can be also produced in the wind from neutron-rich regions of the disk, though the estimated yields have large uncertainties. } ", "introduction": "Observational evidences have been accumulated for a connection between gamma-ray bursts (GRBs) and supernovae (SNe): association of SN 1998bw and GRB 980425\\cite{Ga98} and SN 2003dh in afterglow of GRB 030329.\\cite{Hj03} A {\\itshape collapsar} model is one of promising scenarios to explain a huge gamma-ray production in GRBs and GRB/SN connections.\\cite{MW99,MWH02} During collapse of massive stars, stellar material greater than several solar masses falls back on a new-born black hole with extremely high accretion rates ($\\le 1 M_\\odot\\,$s$^{-1}$).\\cite{WW95} An accretion disk forms around the hole due to the angular momentum of the fallback material.\\cite{MW99,MWH02,MNHNS97} In the context of the collapsar model, jet-like explosion driven by neutrino annihilation and nucleosynthesis in the jet has been investigated.\\cite{NMYTS03} Although \\nuc{Ni}{56} with high velocity ($> 0.1 c$) can be massively produced , it is not sufficient for an observed amount in SN 1998bw.\\cite{NMYTS03} In addition to the production via the jet, massive synthesis of \\nuc{Ni}{56} is also suggested in winds launched from the accretion disk.\\cite{MW99,PWH03} Neutron-rich nuclei may be produced through r-process inside the wind ejected from an inner, neutron-rich region of the disk.\\cite{PWH03} In the present paper, we examine nucleosynthesis inside a GRB accretion disk and investigate abundance change through the wind launched from the disk. ", "conclusions": "We have investigated nucleosynthesis inside the accretion disk associated with GRBs and inside winds launched from an inner region of the disk using the one-dimensional disk and wind models and the nuclear reaction network. The initial composition of accreting gas is taken to be that of an O-rich layer of a 20 $M_{\\odot}$ star before the core collapse. We have found that the disk consists of five layers characterized by dominant elements: \\nuc{O}{16}, \\nuc{Si}{28}, \\nuc{Fe}{54} (and \\nuc{Ni}{56}), \\nuc{He}{4}, and nucleons, and the individual layers shift inward with keeping the overall profiles of compositions as the accretion rate decreases. \\nuc{Ni}{56} are abundantly ejected through the wind from the Fe-rich, He-rich and nucleon-rich disk layers with the electron fraction $\\simeq 0.5$. In addition to iron group elements, heavier elements than Cu, in particular \\nuc{Cu}{63} and \\nuc{Zn}{64}, are massively produced via the wind. Various neutron-rich nuclei can be produced through the wind from neutron-rich regions of the disk in our simple wind model, though the estimated yields have large uncertainties." }, "0405/astro-ph0405383_arXiv.txt": { "abstract": "Polarizability tensor of a strongly magnetized plasma and the polarization vectors and opacities of normal electromagnetic waves are studied for the conditions typical of neutron star atmospheres, taking account of partial ionization effects. Vacuum polarization is also included using a new set of fitting formulae that are accurate for wide range of field strengths. The full account of the coupling of the quantum mechanical structure of the atoms to their center-of-mass motion across the magnetic field is shown to be crucial for the correct evaluation of the polarization properties and opacities of the plasma. The self-consistent treatment of the polarizability and absorption coefficients proves to be necessary if the ionization degree of the plasma is low, which can occur in the atmospheres of cool or ultramagnetized neutron stars. Atmosphere models and spectra are presented to illustrate the importance of such self-consistent treatment. ", "introduction": "\\label{sect:intro} In recent years, thermal or thermal-like radiation has been detected from several classes of isolated neutron stars (NSs): radio pulsars with typical magnetic fields $B\\sim10^{12}$--$10^{13}$ G, ``dim'' NSs whose magnetic fields are mostly unknown, anomalous X-ray pulsars and soft gamma-ray repeaters with $B$ possibly $\\sim10^{14}$--$10^{15}$ G \\citep*[see, e.g.,][ for reviews]{BeckerAschenbach,Haberl-COSPAR,IMS02,PZ03}. The spectrum of thermal radiation is formed in a thin atmospheric layer (with scale height $\\sim0.1$--$10$ cm and density $\\rho\\sim10^{-2}$--$10^3$ \\gcc) that covers the stellar surface. Therefore, a proper interpretation of the observations of NS surface emission requires understanding of radiative properties of these magnetized atmospheres. \\citet*{Shib92} \\citep*[see also][]{SZ95,Pavlov95} presented the first model of the NS atmospheres with strong magnetic fields, assuming full ionization. Variants of this model were constructed by \\citet*{Zane00,Zane01,HoLai,HoLai02,Ozel01,Lloyd03}. An inaccurate treatment of the absorption due to free-free transitions in strong magnetic fields in the earlier models \\citep{Pavlov95} has been corrected by \\citet{PC03}; this correction has been taken into account in later models \\citep{hoetal03,Ho-COSPAR,Lloyd03}. Recent work \\citep{HoLai02,LaiHo02,LaiHo03a} has shown that in the magnetar field regime ($B\\gtrsim 10^{14}$~G) vacuum polarization significantly affects the emergent spectrum from the atmosphere; for weaker fields, vacuum polarization can still leave an unique imprint on the X-ray polarization signals \\citep{LaiHo03b}. Because the strong magnetic field significantly increases the binding energies of atoms, molecules, and other bound states \\citep[see][ for a review]{Lai01}, these bound states may have abundances appreciable enough to contribute to the opacity in the atmosphere. For calculation of this contribution, the non-trivial coupling of the center-of-mass (CM) motion of the atom to its internal structure \\citep[e.g.,][ and references therein]{P94} can be important. Also, because of the relatively high atmosphere density, a proper treatment should take account of the plasma nonideality that leads to Stark broadening and pressure ionization. Recently, thermodynamically consistent equation of state (EOS) and opacities have been obtained for a magnetized, partially ionized H plasma for $8\\times10^{11}$ G $\\la B\\leq10^{15}$ G \\citep*{PCS99,PC03,PC04}. These EOS and opacities have been implemented by \\citet{hoetal03,Ho-COSPAR} for modeling NS atmospheres. For the typical field strengths $B=10^{12}$--$10^{13}$~G this modeling showed that, although the spectral features due to neutral atoms lie at extreme UV and very soft X-ray energy bands and therefore are difficult to observe, the continuum flux is also different from the fully ionized case, especially at lower energies, which can affect fitting of the observed spectra. For the superstrong field $B\\gtrsim 10^{14}$~G, \\citet{hoetal03} showed that the vacuum polarization effect not only suppresses the proton cyclotron line, but also suppresses spectral features due to bound species. It is well known \\citep[e.g.,][]{Ginzburg,Mesz} that under typical conditions (e.g., far from the resonances) radiation propagates in a magnetized plasma in the form of two so-called extraordinary and ordinary normal modes. The polarization vectors of these modes, $\\eX$ and $\\eO$ are determined by the Hermitian part ($\\ChiH$) of the complex polarizability tensor ($\\bm{\\chi}$) of the plasma. Our previous treatment of these modes in partially ionized atmospheres \\citep{PC03,PC04,hoetal03,Ho-COSPAR} was not quite self-consistent, because the effect of the presence of the bound states on the polarization of normal modes was neglected: we adopted the same polarization vectors as in the fully ionized plasma, assuming \\citep{PC03} that the effect of bound states on these vectors should be small provided the ionization degree of the plasma is high. However, this hypothesis (related to $\\ChiH$) may be called into question, based on the observation that the absorption coefficients (corresponding to the {anti-Hermitian} part of the complex polarizability tensor, $i\\ChiA$) are strikingly affected by the presence of even a few percent of the atoms. In this paper, we study the polarizability tensor, the polarization vectors of the normal waves, and the opacities of the partially ionized nonideal hydrogen plasma in strong magnetic fields in a self-consistent manner, using the technique applied previously by \\citet{BulikPavlov} to the case of a monatomic ideal hydrogen gas. In \\S\\ref{sect:general} we introduce basic definitions and formulae to be used for calculation of the plasma polarizability (new fitting formulae for the vacuum polarizability are given in the Appendix). An approximate model based on a perturbation theory, which explains the importance of the CM coupling for plasma polarizability, is described in \\S\\ref{sect:pert}. The results of numerical calculations of the plasma polarizability are presented in \\S\\ref{sect:chi-real}, and the consequences for the polarization and opacities of the normal modes are discussed in \\S\\ref{sect:res}. In \\S\\ref{sect:spectra} we present examples of NS thermal spectra, calculated using the new opacities, compared with the earlier results. In \\S\\ref{sect:concl} we summarize our results, outline the range of their applicability, and discuss unsolved problems. ", "conclusions": "\\label{sect:concl} We have studied the polarizability and electromagnetic polarization modes in a partially ionized, strongly magnetized hydrogen plasma. The full account of the coupling of the quantum mechanical structure of the atoms to their center-of-mass motion across the magnetic field is shown to be crucial for the correct evaluation of the polarization properties and opacities of the plasma. The self-consistent treatment of the polarizability and absorption coefficients is ensured by use of the Kramers-Kronig relation. Such treatment proves to be important if the ionization fraction of the plasma is low ($\\lesssim50$\\%). For high degree of ionization ($\\gtrsim 80$\\%), the polarizability of a fully ionized plasma remains a good approximation, just as previously assumed \\citep{PC03}. This approximation was adopted in the NS atmosphere models built in \\citet{hoetal03,Ho-COSPAR}. A comparison with updated spectra based on the self-consistent treatment (\\S\\ref{sect:spectra}) shows that this approximation is satisfactory if $B\\lesssim10^{13}$~G and $\\Teff\\gtrsim10^6$~K. The self-consistent treatment is needed in the atmospheres of cool or ultramagnetized NSs, with relatively low degrees of ionization. \\begin{figure}\\epsscale{1.} \\plotone{f131.eps} \\caption{Same as in Fig.~\\protect\\ref{fig:f121} but for higher field strength, $B=10^{13}$~G. \\label{fig:f131}} \\end{figure} There are several limitations of the present model, which may become important for the magnetar fields and/or for low $\\Teff$. While H atoms are treated accurately in our calculations of the EOS and opacities, H$_2$ molecules are included in the EOS using the static approximation (i.e., without their CM coupling) and neglected in the opacities. Other bound species, such as H$_2$$^+$ \\citep[e.g.,][]{TurbLop03}, H$_3$$^{2+}$ \\citep{LopTurb}, and H$_n$ chains \\citep*{LSS92} are not included. Moreover, the NS may have a condensed surface, with negligible vapor above it (\\citealt{LS97}; \\citealt{Lai01}). For a NS with mass $M=1.4~\\Msun$ and radius $R=10$ km, estimates at $\\log B$ (G)$=13.5$--15 \\citep[based on][]{PC04} indicate that a thick H atmosphere will be present, and a condensed surface will not occur, provided $\\Teff \\gtrsim 3.8\\times10^5\\,(B/10^{14}\\mbox{ G})^{1/4}$~K; slightly higher $\\Teff$ is needed to ensure negligible abundance of molecules. Another uncertainty in ultramagnetized NS atmospheres is the dense plasma effect: the decoupling layer for photons in the atmosphere (where optical depth $\\approx1$) may occur at high density where the electron plasma frequency exceeds the photon frequency \\citep[e.g.,][]{hoetal03,Lloyd03}. The present treatment is not applicable for such cases. In addition, construction of reliable atmosphere models at $B\\gtrsim 10^{14}$~G requires solution of the problem of mode conversion \\citep{LaiHo03a}. Furthermore, for fitting observed spectra one should construct a grid of models with different field orientations and a range of field strengths, and produce angle- and field-integrated synthetic spectra for an assumed field geometry. Since all the discussed spectral resonances are $B$-dependent, and some of them are $\\theta_B$-dependent, we expect that such integration will somewhat smooth the spectral features \\citep{HoLai04}." }, "0405/astro-ph0405456_arXiv.txt": { "abstract": "{We derive in this Letter the SZ effect induced by the secondary electrons produced in the annihilation of Weakly Interacting Massive Particles (assumed here to be neutralinos) in gravitationally bound structures dominated by Cold Dark Matter (CDM). We show that the DM induced SZ effect has a specific spectral shape and an amplitude which increases for decreasing neutralino mass $M_{\\chi}$. The available SZ data on the Coma cluster set an upper limit on the quantity $\\langle \\sigma V \\rangle_A n^{2}_{\\chi}$ which can be combined with the WMAP constraints on $\\Omega_m h^2$ to restrict the available neutralino models in the $\\langle \\sigma V \\rangle_A - M_{\\chi}$ plane. We delineate various potential applications of this method to constraint the physical properties of the Dark Matter particles from the study of galaxy clusters and dwarf spheroidal galaxies in the light of the next coming high-sensitivity SZ experiments. ", "introduction": "Dark Matter (DM) annihilation in the halo of galaxies and galaxy clusters have relevant astrophysical implications. In fact, if DM is constituted by weakly interacting massive particles (WIMPs), their annihilation can produce direct and indirect signals such as observable fluxes of positrons (e.g., Silk \\& Srednicki 1984, Kamionkowski \\& Turner 1991, Baltz \\& Edsjo 1999) antiprotons (e.g., Bottino et al. 1998), gamma rays (e.g., Chardonnet et al. 1995, Colafrancesco \\& Mele 2001), neutrinos (e.g., Gondolo \\& Silk 1999, Hooper \\& Silk 2004), radio emission (e.g., Colafrancesco \\& Mele 2001), heating of the hot intra-cluster gas (Totani 2004, Colafrancesco 2004). We do not have yet however, at present, a definite detection of these emission features originating from DM annihilation. As an alternative strategy, we explore here the consequences of the Compton scattering between the secondary electrons produced from the WIMP anihilation in massive DM halos, like galaxy clusters, and the CMB photon field. Galaxy clusters are gravitationally dominated by Cold Dark Matter for which the leading candidate is the lightest supersymmetric (SUSY) particle, plausibly the neutralino $\\chi$. Experimental and theoretical considerations for having a cosmologically relevant neutralino DM lead to bound its mass $M_{\\chi}$ in the range between a few GeV to a few hundreds of GeV (e.g., Bottino et al. 2003, Belanger 2003). The decays of neutralino annihilation products (fermions, bosons, etc.) yield, among other particles, energetic electrons and positrons up to energies comparable to the neutralino mass. Here we notice that these energetic electrons and positrons (hereafter we will refer to these particles as electrons because their distinction is not essential for our purpouses) can interact with the CMB photons and up-scatter them to higher frequencies producing a peculiar Sunyaev-Zel'dovich (1972, 1980, hereafter SZ) effect with specific spectral and spatial features. In this Letter, we will describe the specific feature of the SZ effect produced by DM annihilation, SZ$_{DM}$, in galaxy clusters and we will discuss the possibility to disentangle such specific SZ$_{DM}$ effect from the other sources of SZ effect which are present in the same structures. We will finally discuss the future experimental prospects for the detection of the SZ$_{DM}$ effect in DM halos. The relevant physical quantities are calculated using $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$ and a flat, vacuum-dominated CDM ($\\Omega_m = 0.3, \\Omega_{\\Lambda}=0.7$) cosmological model. ", "conclusions": "We have shown that the SZ effect induced by secondary electrons produced in $\\chi \\chi$ annihilation is an unavoidable consequence of the presence and of the nature of Dark Matter in large-scale structures. The analysis of the DM induced SZ effect in galaxy clusters provides a complementary probe for the presence and for the nature of DM in cosmic structures. The available SZ observations on the Coma cluster (see Fig.\\ref{fig.sigmavmchilimits}) can already set a lower limit to the neutralino mass of $M_{\\chi} \\simgt 17 - 20$ GeV ($M_{\\chi} \\simgt 13$ GeV at 90 \\% c.l. with the adopted value of $\\langle \\sigma V \\rangle_A$), which are consistent with the limits set by accelerators (e.g., Belanger et al. 2003). The SZ$_{DM}$ signal does not strongly depend on the assumed DM density profile at intermediate angular distances from the cluster center and on the DM clumpiness since $y_{DM}$ is the integral of the total $P_{DM}$ along the line of sight. The presence of a substantial SZ$_{DM}$ effect is likely to dominate the overall SZ signal at frequencies $x\\simgt 3.8-4.5$ providing a negative total SZ effect (see Fig.\\ref{fig.sz_coma_tot_mchi102030}). It is, however, necessary to stress that in such frequency range there are other possible contributions to the SZ effect, like the kinematic effect and the non-thermal effect which could provide additional biases (see, e.g., Colafrancesco et al. 2003). Nonetheless, the peculiar spectral shape of the $SZ_{DM}$ effect is quite different from that of the kinematic SZ effect and of the thermal SZ effect and this result allows to disentangle it from the overall SZ signal. An appropriate multifrequency analysis of the overall SZ effect based on observations performed on a wide spectral range (from the radio to the sub-mm region) is required, in principle, to separate the various SZ contributions and to provide an estimate of the DM induced SZ effect. In fact, simultaneous SZ observations at $\\sim 150$ GHz (where the SZ$_{DM}$ deepens the minimum with respect to the dominant thermal SZ effect), at $\\sim 220$ GHz (where the SZ$_{DM}$ dominates the overall SZ effect and produces a negative signal instead of the expected $\\approx$ null signal) and at $\\simgt 250$ GHz (where the still negative SZ$_{DM}$ decreases the overall SZ effect with respect to the dominant thermal SZ effect) coupled with X-ray observations which determine the gas distribution within the cluster (and hence the associated dominant thermal SZ effect) can separate the SZ$_{DM}$ from the overall SZ signal, and consequently, set constraints on the neutralino mass. Observations of the radio-halo emission in the cluster can provide an estimate of the cosmic-ray electron population and consequently an estimate of the associated non-thermal SZ effect (which is usually quite small and with a different spectral shape at high frequencies, see e.g., Colafrancesco et al. 2003). The high sensitivity planned for the future SZ experiments, especially at frequencies $x \\approx 2.5$ and $x \\simgt 3.8$, where the SZ$_{DM}$ more clearly manifests itself, can provide much stringent limits to the additional SZ effect induced by DM annihilation. In this context, the next coming PLANCK-HFI experiment has enough sensitivity to probe in details the contributions of various SZ effects in the frequency range $x \\approx 2 - 5$. Because the amplitude of the SZ$_{DM}$ effect increases with decreasing values of $M_{\\chi}$, the high-sensitivity SZ experiments have - hence - the possibility to set reliable constraints to the nature, amount and spatial distribution of DM in galaxy clusters. We will present elsewhere (Colafrancesco 2004, in preparation) a more extended analysis of DM models in the context of SZ observations. An exciting possibility in this context could be offered by nearby (with small or zero peculiar velocity) systems which are gravitationally dominated by Dark Matter, which contain little or no gas (in either hot or warm forms) and show absence of non-thermal phenomena connected with the presence of cosmic rays. In such ideal DM systems, the major source of SZ effect would be just the one due to the annihilation of the WIMPs. Systems which could be assimilable to the ideal \"pure\" DM halos are dwarf spheroidal galaxies and/or low surface brightness galaxies. These systems seem to be ideal sites for studying the DM annihilation indirect signals which reveal themselves in a variety of astrophysical phenomena, whose main imprint is the specific gamma-ray emission (see, e.g., Evans et al. 2003). In such a context, the possible detection of the DM induced SZ effect will provide an important complementary approach which can be studied by more traditional astronomical techniques." }, "0405/astro-ph0405330_arXiv.txt": { "abstract": "A key project of the Antarctic Submillimeter Telescope and Remote Observatory reported by \\citet{martin04} is the mapping of CO $J=4\\rightarrow3$ and $J=7\\rightarrow6$ \\, emission from the inner Milky Way, allowing determination of gas density and temperature. Galactic center gas that \\citet{binney91} identify as being on $x_2$ orbits has a density near $10^{3.5} \\, \\mathrm{cm ^{-3}}$, which renders it only marginally stable against gravitational coagulation into a few Giant Molecular Clouds, as discussed by \\citet{elmegreen94}. This suggests a relaxation oscillator mechanism for starbursts in the Milky Way, where inflowing gas accumulates in a ring at 150 pc radius for approximately 20 million years, until the critical density is reached, and the resulting instability leads to the sudden formation of giant clouds and the deposition of $4 \\times 10^7 {\\mathrm{M_{\\sun}}}$ of gas onto the Galactic center. ", "introduction": "Dynamics of gas in the inner few kiloparsecs of the Milky Way are dominated by the non-axisymmetric gravitational potential of the central bar. The properties of this bar are now well known, and there is good agreement between observations of gas motion and model fits to the potential \\citep{jenkins94,gerhard99,hafner00,bissantz03}. As suggested by \\citet{binney91}, the gas tends to be found on families of closed orbits which are not self-intersecting. All non-closed orbits and some closed orbits are self-intersecting. Gas on such orbits will shock and lose energy where the gas streamlines intersect, and the gas will then move inwards to a lower energy orbit. If the gas can find its way onto a family of non-self-intersecting closed orbits, the energy dissipation slows and the timescale for orbital changes lengthens out. \\citet{contopoulos77} described two families of closed orbits in barred galaxies: the ``$x_1$'', which are elongated along the bar and found outside the inner Lindblad resonance (ILR); and the ``$x_2$'', which are more round and can be found near the ILR and inside it. The ILR is located where the epicyclic frequency of a particle orbiting in the Galactic potential resonates with the pattern speed of the bar. This occurs at a radius of approximately 450 pc from the center of the Milky Way \\citep{bissantz03}. Gas which is several kiloparsecs away from the Galactic center tends to be driven inwards until it reaches a region within a few hundred parsecs of the ILR, because the interaction of the bar potential with the gas exerts a negative torque, resulting in loss of angular momentum by the gas \\citep{lyndenbell72,athanassoula88,jenkins94}. Near the ILR this effect disappears, because the net torque there is small or zero, and inwards of the ILR it may even reverse and become positive, so that gas inside the ILR could be driven outwards \\citep{combes88}. Gas therefore accumulates in a ring near the ILR. Unlike the gas further out, the dynamics of this gas depends critically on its self gravity \\citep{elmegreen94,jenkins94}, and therefore on its thermodynamic properties, density in particular. The thermodynamic properties of the gas can be determined by millimeter- and sub\\-mil\\-li\\-meter-wave spectral line observations. The distribution of molecular gas near the ILR is known from extensive surveys in CO and $\\thco$ $J=1\\rightarrow0$ and $J=2\\rightarrow1$ \\citep{bally88a,bitran97,oka98}; these spectral lines show the presence of molecular gas. These lines alone do not, however, determine its density or excitation temperature. Observations of the mid-$J$ lines of CO provide the missing information. Since the low-$J$ states of CO are in local thermodynamic equilibrium (LTE) in almost all molecular gas \\citep{goldreich74}, measurements of mid-$J$ states are critical to achieving a solution of the radiative transfer by breaking the degeneracy between beam filling factor and excitation temperature. A new survey \\citep{martin04} by the Antarctic Submillimeter Telescope and Remote Observatory \\citep[AST/RO,][]{stark01} adds the $J=4\\rightarrow3$ (461 GHz) and $J=7\\rightarrow6$ (807 GHz) rotational lines of CO to existing lower-frequency data \\citep{bally88a}. These data are available on the AST/RO website\\footnote{\\tt http://cfa-www.harvard.edu/ASTRO} for general use. These measurements have recently been modeled using the large velocity gradient (LVG) approximation to determine the gas density and temperature. In this {\\em Letter}, we discuss the implications of the \\citet{martin04} density estimates for Galactic center gas. We apply our new data, specific to the Milky Way, to the general analysis of stability of dense gas near ILR regions in galaxies by \\citet{elmegreen94}. We find that the gas near 150 pc radius is marginally unstable. This suggests that in the past there has been a period of stability and gas build-up. In the future, the instability will create a few giant clouds, resulting in a starburst and the deposition of tens of millions of solar masses of material on the Galactic center. This process repeats with a cycle time determined by the rate at which gas precipitates on the Galactic center region from outside, resulting in starbursts at intervals of approximately 20 million years. ", "conclusions": "" }, "0405/astro-ph0405106_arXiv.txt": { "abstract": "{Stochastic variability in two out of four XMM-Newton observations of XMMU\\thinspace J004303+4115 along with its power spectra and X-ray luminosities suggest a low-mass X-ray binary (LMXB) with a black hole primary. However, Chandra observations resolve the object into two point sources. We use data from 35 Chandra observations to analyse the contributions of each source, and attribute the variability to \\object{CXOM31\\thinspace J004303.2+411528} { (known as r2-3)}, which varies in intensity by a factor of $\\sim$100 between observations. We assume that the power density spectra of LMXBs are governed by the luminosity, and that the transition between types of power density spectra occurs at some critical luminosity { in Eddington units, $l_{\\rm c}$}, that applies to all LMXBs. We use results from these XMM-Newton observations and past results from the available literature to estimate this transition luminosity, and find that all results are consistent with $l_{\\rm c}$ $\\sim$0.1 in the 0.3--10 keV band. CXOM31\\thinspace J004303.2+411528 exhibits a low accretion rate power density spectrum at a 0.3--10 keV luminosity of 5.3$\\pm0.6$$\\times$10$^{37}$ erg s$^{-1}$. Known stellar mass black holes have masses of 4--15 M$_{\\odot}$; hence our observations of CXOM31\\thinspace J004303.2+411528 are consistent with $l_{\\rm c}$ $\\sim$0.1 if it has a black hole primary. ", "introduction": "\\label{intro} \\begin{table*}[!t] \\centering \\caption{Journal of XMM-Newton observations of the \\object{M31} core}\\label{journ} \\begin{tabular}{lllllll} \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} Observation & Date & MJD& Exp & Filter\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} x1 & 25/06/00 (rev0100)& 51720& 34 ks& Medium \\\\ x2 & 27/12/00 (rev0193)& 51906 & 13 ks& Medium\\\\ x3 & 29/06/01 (rev0285)& 52089& 56 ks &Medium & \\\\ x4 & 06/01/02 (rev0381)& 52280 & 61 ks& Thin\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{table*} The Andromeda Galaxy (M31) is an attractive and important target for X-ray astronomy, since it is the nearest spiral galaxy \\citep[760 kpc, ][]{vdb00}, and its X-ray emission is dominated by point sources. These point sources are thought to be mostly X-ray binaries, along with a few foreground objects, background active galactic nuclei (AGN), and supernova remnants (SNR). The two most recent X-ray observatories, Chandra and XMM-Newton, have finally allowed studies of variability in extra-galactic X-ray sources over time-scales of a few hundred seconds, as well as between successive observations. Such short-term time variability is often characteristic of well-studied phenomena and sometimes allows classification of the objects from X-ray observations alone, in conjunction with their X-ray spectral properties{; for example, thermonuclear X-ray bursts \\citep{lvv95} identify an X-ray source as an { X-ray binary} with a neutron star primary}. To date, analysis of the XMM-Newton observations of the core of M31 has resulted in the discovery of a pulsating supersoft source with a period of 865 seconds \\citep{osb01}, the periodic dipping of the X-ray counterpart to the globular cluster Bo\\thinspace 158 \\citep{tru02}, a persistently bright black hole binary \\citep[][ Paper 1]{bok03}, and a Z-source \\citep[][ Paper 2]{bko03}. Meanwhile \\citet{kaa02} identified variability of three sources in a 47 ks Chandra HRC observation of M31, including the black hole binary later identified in Paper 1. The XMM-Newton and Chandra missions { complement} each other well; Chandra provides imaging with exceptional spatial resolution, while XMM-Newton is the most sensitive imaging X-ray observatory yet flown. The current work exemplifies how XMM-Newton and Chandra results can be used together to get more detailed information than is possible from either data set alone. XMMU\\thinspace J004303+4115 appears as a point source when observed with XMM-Newton, but is resolved by Chandra into two objects, 6\\arcsec~apart. \\citet{K02} associate CXOM31\\thinspace J004202.9+411523, the southern source, with { the globular cluster Bo 146}, and report transient behaviour in CXOM31\\thinspace J004303.2+411528, the northern source. { Following \\citet{K02}, we designate the northern source r2-3 and the southern source r2-4}. We find that in XMM-Newton observations, XMMU\\thinspace J004303+4115 exhibits { power density spectra such as are seen in low accretion rate low-mass X-ray binaries (LMXBs), yet at 0.3--10 keV luminosities of 3--12$\\times$10$^{37}$ erg s$^{-1}$ (Sect. 3).} { \\citet[][ hereafter referred to as vdK94]{vdk94} showed that the power density spectra (PDS) of LMXBs with neutron star or black hole primaries are strikingly similar. At low accretion rates, the PDS of LMXBs have almost identical shapes (approximately broken power laws with the spectral index, $\\gamma$, changing from $\\sim$0 to $\\sim$ 1 at frequencies higher than 0.01--1 Hz) and fractional rms amplitudes of a few times 10\\% (vdK94); we shall refer to these as Type A PDS. At higher accretion rates, LMXBs are considerably less variable, with fractional rms amplitudes of only a few percent, and their PDS are described by power laws with $\\gamma$ $\\sim$1--1.5 (vdK94); we will refer to these as Type B PDS. Furthermore, vdK94 proposed that the transition between Type A and Type B PDS occurs at a critical fraction of the Eddington limit. Following vdK94, we define the critical luminosity fraction, $l_{\\rm c}$, as \\begin{equation} l_{\\rm c} = \\frac{L_{\\rm c}}{L_{\\rm Edd}}, \\end{equation} where $L_{\\rm c}$ is the luminosity of transition between Type A and Type B PDS, and $L_{\\rm Edd}$ is the Eddington luminosity. } In Sect. 3 we analyse longterm lightcurves of r2-3 and r2-4 from 35 Chandra observations, paying particular attention to those that were made within 30 days of one of the XMM-Newton observations. If we can associate { Type A} PDS with either source when its luminosity, $L$, exceeds $L_{\\rm c}$ for a neutron star, we can establish that the primary is a black hole. The most likely candidate for such black hole behaviour is r2-3, since most black hole binaries are transients { \\citep[e.g.][]{mr03}}, and globular cluster X-ray sources mostly contain { 1.4 M$_{\\odot}$} neutron stars \\citep[see ][ and references within]{hgl03}. We present in Sect. 3 evidence that r2-3 contains a black hole primary. In Sect. 4 we first establish that r2-3 is located in M31. We then obtain an empirical value for { $l_{\\rm c}$}, using results from these XMM-Newton observations of globular cluster sources in M31 and published results from analysis of a Galactic neutron star LMXB and a Galactic black hole LMXB. { We then use our value of { $l_{\\rm c}$} to calculate $L_{\\rm c}$ for a neutron star with a mass of 3.1 M$_{\\odot}$, the theoretical maximum \\citep{kk78}. We assert that r2-3 exhibits a Type A PDS at a luminosity that exceeds this limit and conclude that the primary in r2-3 is a black hole, consistent { with} its transient behaviour. } ", "conclusions": "In diagnosing the nature of r2-3 we first need to be certain that it is indeed within the M31 galaxy. We applied a column density of 1.0$\\times$10$^{21}$ cm$^{-2}$ and a spectral index of 1.78 to the HRC-I observations that yielded the maximum and minimum ACIS-I { equivalent} count rates of r2-3 and obtained a flux range of 0.012--1.4$\\times$10$^{-12}$ erg cm$^{-2}$ s$^{-1}$ in the 0.3--10 keV band using {\\sc pimms}. This corresponds to a luminosity range of 0.08--9.5$\\times$10$^{37}$ erg s$^{-1}$ for a location in M31. If it was local, it would be most likely to be within 1 kpc, i.e. $<$3 times the scale-height of the Galactic disc, since M31 is 21.6$\\degr$ out of the Galactic plane and 120$\\degr$ from the Galactic Centre. In this situation, the luminosity would be $>$6 orders of magnitude smaller, and hence up to 2 orders of magnitude fainter than the faintest known persistent Galactic LMXB \\citep{wil03}. It would also be up to a factor of $\\sim$3 fainter than the faintest black hole X-ray transient in quiescence \\citep[][ and references within]{tom03}. It is likely that if r2-3 { were} local, it would have an optical counterpart; known absolute V magnitudes (M$_{\\rm V}$) of persistent Galactic LMXBs range from $-$2.5 to 5.6 \\citep{vm95}, and the spectral types of the secondary stars in most black hole X-ray transients have been obtained in quiescence \\citep{cc03}, with { M$_{\\rm V}$} $\\sim$0--9.7. Hence, { if r2-3 were local, we would expect to see an optical counterpart with m$_{\\rm V}$ $\\la$ 19}. However, the nearest optical source in the HST catalogue of \\citet{hai94} is 79$\\arcsec$ away, with a { m$_{\\rm v}$ = 20.6}; hence any optical counterpart to r2-3 would have to be fainter than this. Thus r2-3 is unlikely to be local. In addition, its PDS rules out the possibility that it is a background AGN, as they exhibit spectral breaks at 10$^{-6}$--10$^{-5}$ Hz \\citep{utt02} rather than at the $\\sim$0.03 Hz seen here. We therefore conclude that r2-3 is located in M31. By showing that we observe a { Type A} PDS from r2-3 at a luminosity that { is} too high for a neutron star LMXB we can classify the primary as a black hole. To do this, we must obtain a value for { $l_{\\rm c}$, and assume that this applies to all LMXBs. Then we must show that r2-3 exhibits a Type A PDS at a luminosity greater than $L_{\\rm c}$ for any neutron star.} In the first instance, we looked at seven X-ray sources associated with globular clusters in { x4}, since they are likely to contain 1.4 M$_{\\odot}$ neutron stars. None of { them exhibited Type A PDS; the luminosity range was } $\\sim$2--10$\\times$10$^{37}$ erg s$^{-1}$ (Paper 1). Also, a { Type B} PDS was exhibited by XMMU\\thinspace J004303+4115 in observation x1 at 3$\\times$10$^{37}$ erg s$^{-1}$ (which we associate with r2-4 alone, see Sect. 3). These results suggest that { $L_{\\rm c}$}$\\la$ 2$\\times$10$^{37}$ erg s$^{-1}$; hence { $l_{\\rm c}$ $\\la$ 0.1, assuming hydrogen accretion onto a 1.4 M$_{\\odot}$ neutron star.} We obtained a vital clue to { $l_{\\rm c}$} from the Galactic neutron star \\object{LMXB 4U\\thinspace 1705$-$44}. \\citet{lang89} analysed data from four EXOSAT observations of 4U\\thinspace 1705$-$44; they found that it exhibited a { Type A PDS in the faintest observation but a Type B PDS in the next faintest}. In their previous analysis of the observations \\citep{lang87}, they obtained 1--11 keV fluxes of 1.3$\\times$10$^{-9}$ and 1.8$\\times$10$^{-9}$ erg { cm$^{-2}$} s$^{-1}$ for these two observations. Hence an accurate distance to 4U\\thinspace 1705$-$44 would yield a tight constraint on { $l_{\\rm c}$}. \\citet{cs97} estimate a distance of 11 kpc, using the most luminous X-ray burst as a standard candle; this constitutes an upper limit to the distance \\citep[see e.g.][ and references within]{kul03}. \\citet{cor03} also estimate the distance to 4U\\thinspace 1705$-$44 using bursts, but give a distance of 8.9 kpc with an assumed uncertainty of 30\\%. Hence { $l_{\\rm c}$} = 0.08$^{+0.08}_{-0.05}$. { Similarly, \\object{GS\\thinspace 2023+338} (\\object{V404 Cygni}) is a Galactic black hole LMXB; the most likely mass for the primary is 12 M$_{\\odot}$ \\citep{shab94}. It was discovered with Ginga during an outburst in 1989 \\citep{mak89} and exhibited Type A PDS at 2--37 keV luminosities $>$3$\\times$10$^{38}$ erg s$^{-1}$ \\citep[][ and references within]{miy92,oost97}. The X-ray spectrum for V404 Cygni was described by a power law with $\\alpha$ = 1.0--1.4 \\citep{miy92}, hence $l_{\\rm c}$ $\\ga$ 0.06 in the 0.3--10 keV band. } These three sets of results are all consistent with { $l_{\\rm c}$ $\\sim$0.10 in the 0.3--10 keV band; this supports the idea of a constant $l_{\\rm c}$ proposed by vdK94}. { Now, r2-3 appears to exhibit a Type A PDS at a 0.3--10 keV luminosity of 5.3$\\pm$0.6$\\times$10$^{37}$ erg s$^{-1}$; this is a factor of $\\sim$3 higher than $L_{\\rm c}$ for a 1.4 M$_{\\odot}$ neutron star. Indeed, for the maximum mass of a neutron star \\citep[i.e. 3.1 M$_{\\odot}$,][]{kk78}, $L_{\\rm c}$ $\\sim$4$\\times$10$^{37}$ erg s$^{-1}$. { Known stellar-mass black holes have masses over the range 4--15 M$_{\\odot}$, hence our results from r2-3 are consistent with $l_{\\rm c}$ = 0.1 if the primary is a black hole. } }" }, "0405/astro-ph0405276_arXiv.txt": { "abstract": "{We present a numerical investigation of the development of Rayleigh-Taylor instability at the interface between an expanding Pulsar Wind Nebula and its surrounding Supernova Remnant. These systems have long been thought to be naturally subject to this kind of instability, given their expansion behavior and the density jump at the contact discontinuity. High resolution images of the Crab Nebula at optical frequencies show the presence of a complex network of line-emitting filaments protruding inside the synchrotron nebula. These structures are interpreted as the observational evidence that Rayleigh-Taylor instability is in fact at work. The development of this instability in the regime appropriate to describe Supernova Remnant-Pulsar Wind Nebula systems is non-trivial. The conditions at the interface are likely close to the stability threshold, and the inclusion of the nebular magnetic field, which might play an important role in stabilizing the system, is essential to the modeling. If Rayleigh-Taylor features can grow efficiently a mixing layer in the outer portion of the nebula might form where most of the supernova material is confined. When a magnetic field close to equipartition is included we find that the interface is stable, and that even a weaker magnetic field affects substantially the growth and shape of the fingers. ", "introduction": "Pulsars are rapidly rotating magnetized neutron stars that usually form as the result of a Supernova (SN) explosion. As a consequence of the electromagnetic torques acting on it, a pulsar releases most of its rotational energy in the form of a relativistic magnetized wind. The wind is usually thought to be made of electron-positron pairs and to carry a magnetic field that far enough from the light cylinder is almost purely toroidal (\\cite{michel99}; \\cite{goldreich69}). This outflow is highly relativistic, with a terminal Lorentz factor in the range $10^4$-$10^7$. Its confinement by the surrounding Supernova Remnant (SNR) generates a nebula of relativistically hot material that shines through synchrotron and Inverse Compton emission from radio wavelengths up to $\\gamma$-rays: this is what we call a pulsar wind nebula (PWN) or ``plerion''. During a SN explosion as much as $10^{51}$ erg of energy are released in the form of a blast wave that produces a strong shock propagating in the surrounding medium. The ejected material is initially heated by the blast wave and set into motion. As the ejecta expand, their thermal pressure finally becomes so low as to be dynamically unimportant: from this moment on the expansion can be approximated as homologous (\\cite{chevalier89}; \\cite{matzner99}). This phase is referred to as ``free expansion'' of the ejecta. The evolution of the PWN inside the free expanding ejecta depends on many different parameters such as the pulsar luminosity, the flux anisotropies of the pulsar wind (\\cite{komissarov03}; \\cite{delzanna04}), the density and velocity distribution in the ejecta (\\cite{dwarkadas98}; \\cite{featherstone01}; \\cite{blondin96}), as well as the presence of large and/or small scale anisotropies (\\cite{chevalier89}; \\cite{campbell03}). As a consequence, the detailed modeling of a single PWN-SNR system requires the knowledge of a number of parameters depending on the specific conditions, that are usually unknown. The simplest approximation one can make for the time evolution of the PWN size is obtained assuming constant pulsar luminosity and spherical symmetry (\\cite{chevalier92}; \\cite{swaluw01}; \\cite{bucciantini03}; \\cite{bucciantini04}). If a radial power law density profile, $\\rho \\propto r^{-\\alpha}t^{\\alpha-3}$, is further assumed for the SN ejecta, then the PWN size evolves as $t^{(6-\\alpha)/(5-\\alpha)}$. For a more detailed description of the various phases of the PWN-SNR evolution see Bucciantini et al. (2003) and references therein. The interface between the synchrotron nebula and the swept up shell of ejecta has been thought to be Rayleigh-Taylor (hereafter RT) unstable (\\cite{chevalier75}; \\cite{bandiera83}). In the case of the Crab Nebula the RT instability is expected to be at the origin of the complex network of emission-line filaments protruding into the PWN (\\cite{hester96}, H96 hereafter). The recent images by H96 show that these ``radial'' filaments are joined at their basis by faint, thin , ``tangential'', and often somewhat arcuate features, which are interpreted as tracing the location of the RT unstable interface. The filamentary structure presents a clear hierarchy as expected from a multimode instability. Simulations of the RT instability in the first phase of the PWN-SNR evolution have been presented by Jun in a classical hydrodynamical (HD) regime (\\cite{jun98}, hereafter J98). However, as H96 pointed out, the standard model for PWNe (\\cite{kennel84}) leads to believe that many of them are magnetically dominated at the contact discontinuity with the SNR. In addition, within the standard framework, the nebular magnetic field, usually expected to be purely toroidal, is tangential to the contact discontinuity. A parallel magnetic field is expected to have a stabilizing effect and possibly even to suppress the formation of fingers (\\cite{chandrasekhar61}; \\cite{wang83}; \\cite{jun95}). In the recent work by H96 a comparison with results from classical magnetohydrodynamical (MHD) simulations of the RT instability (\\cite{jun95}) was carried out. The results suggest that, in the case of the Crab Nebula, the magnetic field should be close to the critical value for stability. However, the conditions at the contact discontinuity between PWNe and SNRs are quite different from those adopted both in the formulation of the standard theory (\\cite{chandrasekhar61}) and in the existent MHD simulations (\\cite{jun95}), as we will discuss in the following. In this paper we present a study of the RT instability in the presence of a tangential magnetic field in the context of PWN-SNR systems. Our analysis is carried out by means of 2D special relativistic MHD (RMHD) simulations. In Section 2 we review the standard theory for plane-parallel RT instability, discussing how the main results are modified in the situation under investigation. In Section 3 the numerical method and the initial conditions for the simulations are described. Section 4 is dedicated to the numerical results both in the HD and MHD regimes. In Section 5 we finally summarize our conclusions. ", "conclusions": "In this paper, for the fist time, we present simulations of RT instability for PWNe expanding into the SNR ejecta, including the effect of a magnetic field at the boundary. Simulations were carried out for various initial angular perturbations. In the HD regime our simulations confirm previous results by J98 concerning the nebular size and fraction of mass contained in the mixing layer. The higher resolution we adopted allows us to follow the detailed evolution of the finger structure. We find that finger fragmentation is likely to happen only once a well developed turbulent mixing layer has formed, and that, for large scale perturbations, TS instability seems to dominate over RT, leading to an overall deformation of the swept-up shell of ejecta. We were able to observe the dragging exerted by secondary KH and the formation of mushroom caps. From our simulations it seems that fingers do not grow enough to penetrate the relativistic wind region and the mixing layer extends to about one quarter of the nebular radius. The introduction of a magnetic field turns out to be a key ingredient for a correct understanding and modeling of the RT instability. We have shown that the stability criterion adapted to the self-similar PWN-SNR evolution, does not depend on pulsar wind luminosity, SN mass and energy and does not change during the nebula evolution. The result is that the ratio between the shell density and the critical density is close to unity if a magnetic field around equipartition is assumed at the PWN boundary. This result is confirmed by our simulations that show that magnetic field close or above equipartition can completely suppress the RT instability even for large scale perturbations. A weaker magnetic field is able to reduce the growth of the finger leading to a round rather than elongated protuberance attached to the ejecta shell. A non-negligible magnetic field can suppress the secondary KH completely: no turbulent mixing layer is formed, and fingers are thicker than in the HD case. This is the main difference with respect to former simulations in the pressure dominated regime, where the standard stability criterion was found to give the correct result. As shown by H96 an efficient cooling is required to explain the RT filamentary network observed in the Crab Nebula. Our simulations however point out that this requirement is even stronger than previously thought, thus favoring the hypothesis that the fingers formed during the early phases of the system evolution, when the density of the shell was higher. An alternative to efficient cooling would be the presence of dense clumps formed during the SN that may resist the expansion of the nebula. Recent work on axisymmetric pulsar winds seems also to suggest a third possible explanation for the formation of RT fingers and possibly their latitude distribution. If turbulent large scale convective cells are formed, as a consequence of ram pressure gradient in the wind, there might be regions at the boundary where the magnetic field is below equipartition and the RT instability criterion is satisfied. This scenario deserves more attention, however 3D global simulations with high resolution are still too much demanding in terms of computational time." }, "0405/astro-ph0405089_arXiv.txt": { "abstract": " ", "introduction": "} In X-ray binary systems a neutron star or a black hole accretes matter from a nearby companion star, either by Roche-lobe overflow or by wind accretion. Due to conservation of angular momentum the accreted matter does not directly fall onto the compact object, but forms an accretion disk around it. In this accretion disk, angular momentum is transported outward as the matter spirals in. A large amount of gravitational energy (up to \\mbox{$\\sim$10$^{38}$ erg s$^{-1}$}) is released when the matter approaches the compact object, heating the inner accretion disk to very high temperatures (\\mbox{$\\sim$10$^7$ K}) and causing it to emit X-rays. X-ray binaries can be divided into high-mass X-ray binaries and low-mass X-ray binaries (LMXBs) after the mass of the companion star. In the first group the mass of the companion star usually exceeds ten solar masses and the accretion is driven by its strong wind, while in LMXBs the companion's mass is below one solar mass and the accretion is driven by Roche-lobe overflow. Systems which contain a donor star with a mass between one and ten solar masses are rare, due to the low efficiency of wind accretion and the instability of the mass transfer through Roche-lobe overflow when the donor is more massive than the receiving star. LMXBs can be further classified into {\\it persistent} and {\\it transient} sources depending on their long-term X-ray variability (see \\S~\\ref{subsection:outburst_phase}). A sub-group of the transient sources (the {\\it quasi-persistent} neutron-star X-ray transients) has recently led to advances in our understanding of the properties of the cores and crusts of neutron stars and it is this research that is the focus of this chapter. \\subsection{Neutron-star X-ray transients: outburst phase \\label{subsection:outburst_phase}} \\begin{figure}[t] \\begin{center} \\begin{tabular}{c} \\psfig{figure=rwijnands_fig1.ps,width=13cm,angle=-90} \\end{tabular} \\caption{ {\\it RXTE}/ASM light curves (since January 1, 1996; i.e., since the start of the {\\it RXTE} mission) of the neutron-star X-ray transients XTE J1806--246 (top), MXB 1730--335 (middle; this source is also known as the Rapid Burster, located in the globular cluster Liller 1), and MXB 1659--29 (bottom). It is clear that a wide variety of outburst durations and frequencies has been observed for different transients. The ASM data points plotted represent four-day averages for XTE J1806--246 and MXB 1730--335 , but seven-day averages for MXB 1659--29. \\label{fig:transients_lc} } \\end{center} \\end{figure} The neutron-star X-ray transients form a special group among neutron-star LMXBs. They are usually very dim, with luminosities of $10^{32-34}$~\\Lunit, but occasionally they exhibit violent outbursts during which their X-ray luminosity increases by several orders of magnitude to $10^{36-38}$ \\Lunit~(e.g., Chen {\\em et al.}~1997). These outbursts typically last for several weeks to months before the systems turn off again. For most of these sources only one outburst has been observed, although several have shown multiple outbursts (see Fig.~\\ref{fig:transients_lc} for typical examples of light curves as obtained with the all sky monitor [ASM] aboard the {\\it Rossi X-ray Timing Explorer} [{\\it RXTE}] satellite). These outbursts are very likely the result of considerable increases in the mass accretion rates onto the neutron stars in these systems, although the exact physical processes behind these outbursts are still not well understood (see, e.g., Lasota 2001 for a review of outburst physics). Among the transients there is a special sub-class of sources which do not turn off after a few weeks or months but which remain active for many years (the 'quasi-persistent' transients). The best examples of quasi-persistent transients harboring neutron stars are EXO 0748--676 and GS 1826--238 (which are still active), and MXB 1659--29 and KS 1731--260 (which turned off recently). In addition, several neutron-star LMXBs which once were thought to be persistent suddenly turned off (e.g., 4U 2129+47, X1732--304; XB 1905+000; Pietsch {\\em et al.}~1986; Guainazzi {\\em et al.}~1999) and they should also be considered quasi-persistent transients (throughout this chapter when we refer to 'quasi-persistent transients' we mean {\\it quasi-persistent neutron-star X-ray transients} unless otherwise noted\\footnote{Note that also quasi-persistent {\\it black-hole} X-ray transients exist, such as GRS 1915+105 (active since May 1992; Paciesas {\\em et al.}~1994) and 4U 1630--47 (active since September 2002; Wijnands {\\em et al.}~2002c), however, we do not discuss these systems in this review.}). During their outbursts, the neutron-star transients (both the ordinary and the quasi-persistent transients) are very similar to persistent neutron-star LMXBs with respect to their X-ray properties and can be readily studied by the X-ray satellites available. However, obtaining high quality X-ray data from the transient systems during their quiescent state still remains a challenge because of the much lower X-ray luminosities. \\subsection{Short duration neutron-star X-ray transients: quiescent phase} Despite the low X-ray luminosities of neutron-star X-ray transient during their quiescent state, they can still be detected with sensitive imaging instruments. Although several intrinsically bright and/or nearby systems were already detected with older generation X-ray satellites (e.g., {\\it EXOSAT}, {\\it ROSAT}, {\\it ASCA}, and {\\it BeppoSAX}; van Paradijs {\\em et al.}~1987; Verbunt {\\em et al.}~1994; Garcia 1994; Asai {\\em et al.}~1996, 1998; Campana {\\em et al.}~1998b, 2000; Garcia \\& Callanan 1999; Stella {\\em et al.}~2000), the launch of the {\\it Chandra} and {\\it XMM-Newton} X-ray satellites with their high sensitivity cameras meant a great leap forward in our ability to detect quiescent systems and to obtain good X-ray spectra (see, e.g., Daigne {\\em et al.}~2002; in 't Zand {\\em et al.}~2001; Rutledge {\\em et al.}~2001a, 2001b; Wijnands {\\em et al.}~2001, 2002b, 2003; Campana {\\em et al.}~2002; Jonker {\\em et al.}~2003, 2004a). \\subsubsection{The observed X-ray properties in quiescence} So far, the majority of quiescent neutron-star transients exhibit X-ray spectra that are dominated by a soft ($<$$1$ keV) component which can be accurately described by a thermal model such as a black-body model or a modified black-body model like the neutron-star atmosphere (NSA) models. In these NSA models it is assumed that the depth from which the observed photons emerge from the neutron-star atmosphere increases significantly with the energy of the photons, due to the strong dependency of the opacities on energy. Therefore, high energy photons emerge from deeper and hotter layers than less energetic photons. For a particular temperature, the emerging X-ray spectrum is thus harder than the one that would result if the neutron star would radiate as a pure black body. If indeed the neutron star emits a NSA-like spectrum, then fitting that spectrum with a black-body would overestimate the effective temperature (by a factor of 2) and therefore underestimate the emitting area (often by an order of magnitude; see, e.g., Zavlin {\\em et al.}~1996 and references therein). The NSA models (those assuming a hydrogen atmosphere and a negligible neutron-star magnetic field strength\\footnote{The neutron stars in X-ray transients are assumed to have magnetic field strengths of only $10^8 - 10^9$ Gauss, which is sufficiently low not to affect the spectra emerging from the neutron star. Thus the appropriate NSA models are those which assume a zero magnetic field strength.}) have recently dominated the spectral fits reported in the literature of quiescent neutron-star X-ray transients. This is because NSA models provide a clear physical explanation for the shape of the emitted quiescent spectrum and they yield radii of the emitting area which are consistent with the theoretically expected radii of neutron stars. In contrast, black-body models typically give radii which are significantly lower than those expected for neutron stars. However, it is important to stress that black-body models provide fits to the data that are equally as satisfying as those of NSA models. Thus, currently, we cannot distinguish between these models observationally. When fitting NSA models to the X-ray spectra of most quiescent neutron-star transients, we observe that they have typically an effective temperature (all effective temperatures in this chapter are for an observer at infinity) of 0.1--0.2 keV and a bolometric luminosity between $10^{32}$ and $10^{34}$ ergs s$^{-1}$. Several systems have been found to exhibit an additional spectral component which dominates the spectrum above a few keV and which can be described by a simple power-law model (e.g., Asai {\\em et al.}~1998; Rutledge {\\em et al.}~2001a, 2001b). This component can contribute up to 50\\% to the 0.5--10 keV quiescent flux of a particular system (e.g., Rutledge {\\em et al.}~2001b), although in other systems it cannot be detected and at most 10\\%--20\\% of the 0.5--10 keV flux might be due to such an additional hard spectral component (e.g., Wijnands {\\em et al.}~2003). Although most systems are dominated by the soft thermal component, two systems are known {\\it not} to follow this general trend. Instead, they are dominated by the hard power-law component (contributing $>$90\\% to the 0.5--10 keV quiescent flux) and no thermal component could be conclusively detected. Campana {\\em et al.}~(2002) found that the accretion-driven millisecond X-ray pulsar and X-ray transient SAX J1808.4--3658 had a quiescent spectrum which was dominated by the hard power-law component. Furthermore, its quiescent luminosity was observed to be $5\\times 10^{31}$ ergs s$^{-1}$, which makes it the (intrinsically) faintest quiescent neutron-star transient currently known. Very recently, Wijnands {\\em et al.}~(2004b) found that, EXO 1745--248 in the globular cluster Terzan 5, is the second system with a quiescent X-ray spectrum dominated by the hard power-law component. Again the thermal component could not be detected and it contributed at most 10\\% to the quiescent 0.5--10 keV flux. Although this resembles SAX J1808.4--3658, the 0.5--10 keV luminosity of EXO 1745--248 was a factor of 40 larger than that observed for SAX J1808.4--3658. Currently, it is not understood why SAX J1808.4--3658 and EXO 1745--248 are different from the majority of quiescent neutron-star X-ray transients. \\subsubsection{Theoretical models for quiescent neutron-star X-ray transients} Several theoretical models have been developed to explain the low quiescent X-ray luminosities and the X-ray spectra observed for neutron-star X-ray transients. For example, the X-rays could be due to the residual accretion of matter onto the neutron-star surface or down to the magnetospheric boundary, or the pulsar emission mechanism might be active (see, e.g., Stella {\\em et al.}~1994; Zampieri et al.~1995; Corbet 1996; Campana {\\em et al.}~1998a; Menou {\\em et al.}~1999; Campana \\& Stella 2000; Menou \\& McClintock 2001). The model currently most often used to explain the soft component is the 'cooling neutron star model'. In this model (e.g., Campana {\\em et al.}~1998a; Brown {\\em et al.}~1998) the radiation emitted below a few keV is thermal emission originating from the neutron star surface. Brown {\\em et al.}~(1998) argued that the neutron star core is heated by the nuclear reactions occurring deep in the crust when the star is accreting. This heat is released as thermal emission during quiescence. If the quiescent emission is dominated by the thermal emission of the cooling neutron star, then the quiescent luminosity should depend on the time averaged (over $10^{4-5}$ years) accretion luminosity of the system (Campana {\\em et al.}~1998a; Brown {\\em et al.}~1998). Thus, the quiescent luminosities observed can be directly compared with luminosities predicted using estimations of the long term accretion history of the sources. The version of the cooling neutron star model presented by Brown et al. (1998) was able to explain, at the time of its publication, the luminosities of most of the systems then detected, although it was found that the neutron-star transient Cen X-4 appeared to be less luminous than this model predicted. This could be due to an overestimation of the time-averaged accretion rate of Cen X-4 or due to the existence of enhanced cooling processes in the core of the neutron star (e.g., due to the direct Urca process, pion condensation, or Cooper-pairing neutrino emission) instead of the standard core cooling processes assumed by Brown {\\em et al.}~(1998). Since the publication of the Brown {\\em et al.}~(1998) paper, {\\it Chandra} and {\\it XMM-Newton} have provided us with high quality data on about a dozen quiescent neutron-star systems. We now know that Cen X-4 is not the only system which is colder than expected on the basis of its time-averaged accretion history and the standard cooling model (e.g., Campana {\\em et al.}~2002; Nowak {\\em et al.}~2002). It seems that two groups of sources exist: those which can be explained by assuming standard core cooling and those which require enhanced core cooling processes. Not all characteristics of the quiescent emission can be fully explained by the cooling neutron star model (either using standard or enhanced core cooling). For example, the neutron-star transients Aql X-1, Cen X-4, and SAX J1748.9--2021 (located in the globular cluster NGC 6440) have shown considerable variability in their quiescent properties on time scales ranging from hundreds of seconds to years (Rutledge {\\em et al.}~2002a; Campana \\& Stella 2003; Campana {\\em et al.}~2004; Cackett {\\em et al.}~2004). This variability cannot easily be explained by a cooling neutron star and extra ingredients need to be added to the cooling model (e.g., Ushomirsky \\& Rutledge 2001; Brown {\\em et al.}~2002) or alternative models must be used to explain the quiescent properties (e.g., Campana \\& Stella 2003; Campana {\\em et al.}~2004). Furthermore, the power-law shaped spectral component which dominates the quiescent spectra above a few keV in several systems cannot be explained by the cooling models. The difficulty is even more dire when attempting to explain the hard, power-law component that dominates the quiescent X-ray spectra of SAX J1808.4--3658 and EXO 1748--248. It is conceivable that this component might be described by one or more of the alternative models mentioned above (i.e., those which assume that the neutron star has a non-negligible magnetic field strength). However, the observational results on this component and our understanding of its nature are very limited. ", "conclusions": "Observations of quasi-persistent neutron-star X-ray transients (i.e., those of KS 1731--260 and MXB 1659--29) in quiescence have demonstrated that this class of transients can be used to constrain the structure of neutron stars. However, in order to fully realize the promise of these observations, more work still needs to be done, especially calculating cooling curves specifically for each of the neutron stars in the various quasi-persistent transients given that these curves depend on the long-term accretion history of the source which is quite different among systems. Furthermore, additional quasi-persistent neutron-star transients should be observed and monitored in their quiescent state preferably within the first year after the end of their last accretion episode. Such observations will allow us to determine whether the behavior of KS 1731--260 and MXB 1659--29 is typical among quasi-persistent neutron-star X-ray transients. Differences might be expected not only because the accretion histories of the various sources will differ significantly from each other but also because the crusts (large heat conductivity vs. low heat conductivity) and cores (standard vs. enhanced core cooling) of the neutron stars can behave differently. All the known persistent neutron-star LMXBs are promising candidates should they ever turn off. As shown above, several sources which were thought to be persistent have turned off suddenly, so it is possible that an additional 'persistent' source might be found instead to be a quasi-persistent X-ray transient. Unfortunately, the likelihood of any one of them turning off is low since they have now been found to be active for over 40 years (since the dawn of X-ray astronomy). More promising are those systems which have been seen to suddenly turn on in the last 20 years and which have stayed on ever since (the most promising candidates are EXO 0748--676, GS 1826--24, and XTE J1759--220). It is likely that these systems may turn off in the future and could then be used in a way similar to KS 1731--260 and MXB 1659--29 to constrain the properties of neutron stars." }, "0405/gr-qc0405007_arXiv.txt": { "abstract": "This paper discusses the so called holographic solution, in short \"holostar\". The holostar is the simplest exact, spherically symmetric solution of the original Einstein field equations with zero cosmological constant, including matter. The interior matter-distribution follows an inverse square law $\\rho = 1 / (8 \\pi r^2)$. The interior principal pressures are $P_r = - \\rho$ and $P_\\perp = 0$, which is the equation of state for a collection of radial strings with string tension $\\mu = -P_r = 1/(8 \\pi r^2)$. The interior space-time is separated from the exterior vacuum space-time by a spherical two-dimensional boundary membrane, consisting out of (tangential) pressure. The membrane has zero mass-energy. Its stress-energy content is equal to the holostar's gravitational mass. The total gravitational mass of a holostar can be determined by a proper integral over the Lorentz-invariant trace of the stress-energy tensor. The holostar exhibits properties similar to the properties of black holes. The exterior space-time of the holostar is identical to that of a Schwarzschild black hole, due to Birkhoff's theorem. The membrane has the same properties (i.e. the same pressure) as the fictitious membrane attributed to a black hole according to the membrane paradigm. This guarantees that the dynamical action of the holostar on the exterior space-time is identical to that of a black hole. The holostar possesses an internal temperature proportional to $1/\\sqrt{r}$ and a surface redshift proportional to $\\sqrt{r}$, from which the Hawking temperature and entropy formula for a spherically symmetric black hole are derived up to a constant factor. The holostar's interior matter-state is singularity-free. It can be interpreted as the most compact spherically symmetric (i.e. radial) arrangement of classical strings: The radially outlayed strings are densely packed, each string occupies exactly one Planck area in its transverse direction. This maximally dense arrangement is the reason why the holostar does not collapse to a singularity and why the number of interior degrees of freedom scales with area. Although the holostar's {\\em total} interior matter state has an overall string equation of state, {\\em part} of the matter can be interpreted in terms of particles. The number of zero rest-mass particles within any concentric region of the holostar's interior is shown to be proportional to the proper area of its boundary, implying that the holostar is compatible with the holographic principle and the Bekenstein entropy-area bound, not only from a string but also from a particle perspective. In contrast to a black hole, the holostar-metric is static throughout the whole space-time. There are no trapped surfaces and no event horizon. Information is not lost. The weak and strong energy conditions are fulfilled everywhere, except for a Planck-size region at the center. Therefore the holostar can serve as an alternative model for a compact self-gravitating object of any conceivable size. The holostar is the prototype of a closed system in thermal equilibrium. Its lifetime is several orders of magnitude higher than its interior relaxation time. The thermodynamic properties of the interior space-time can be derived exclusively from the geometry. The local entropy-density $s$ in the interior space-time is exactly equal to the proper geodesic acceleration of a stationary observer, $s = g / \\hbar$. It is related to the local energy-density via $s T = \\rho$. The free energy in the interior space-time is minimized to zero, i.e. $F = E - ST = 0$. Disregarding the slow process of Hawking evaporation, energy and entropy are conserved locally and globally. Although the holostar is a static solution, it behaves dynamically with respect to the interior motion of its constituent particles. Geodesic motion of massive particles in a large holostar is nearly radial and exhibits properties very similar to what is found in the observable universe: Any material observer moving geodesically outward will observe an isotropic outward directed Hubble-flow of massive particles from his local frame of reference. An inward moving observer experiences an inward directed Hubble-flow. The outward motion is associated with an increase in entropy, the inward motion with a decrease. The radial motion of the geodesically moving observer is accelerated, with the proper acceleration falling off over time (for an outward moving observer). The acceleration is due to the interior metric, there is no cosmological constant. Geodesic motion of massive particles is highly relativistic, as viewed from the stationary coordinate system. The $\\gamma$-factor of an inward or outward moving observer is given by $\\gamma \\approx \\sqrt{r/r_0}$, where $r_0$ is a fundamental length parameter which can be determined experimentally and theoretically. $r \\approx 2 r_{Pl}$. Inmoving and outmoving matter is essentially decoupled, as the collision cross-sections of ordinary matter must be divided by $\\gamma^2$, which evaluates to $\\approx 10^{60}$ at the radial position of an observer corresponding to the current radius of the universe. The total matter-density $\\rho$, viewed from the extended Lorentz-frame of a geodesically moving observer, decreases over proper time $\\tau$ with $\\rho \\propto 1 / \\tau^2$. The radial coordinate position $r$ of the observer evolves proportional to $\\tau$. The local Hubble value is given by $H = 1/ \\tau$. Although the relation $\\rho \\propto 1/r^2$ seems to imply a highly non-homogeneous matter-distribution, the large-scale matter-distribution seen by a geodesically moving observer within his observable local Hubble-volume is homogeneous by all practical purposes. The large-scale matter-density within the Hubble-volume differs at most by $\\delta \\rho / \\rho \\approx 10^{-60}$ at radial position $r \\approx 10^{61}$, corresponding to the current Hubble-radius of the universe. The high degree of homogeneity in the frame of the co-moving observer arises from the combined effect of radial ruler distance shrinkage due the radial metric coefficient ($\\sqrt{g_{rr}} \\propto \\sqrt{r/r_{Pl}}$ and Lorentz-contraction in the radial direction because of the highly relativistic geodesic motion, which is nearly radial with $\\gamma \\propto \\sqrt{r/r_{Pl}}$ . The geodesically moving observer is immersed in a bath of zero rest-mass particles (photons), whose temperature decreases with $T \\propto 1 / \\sqrt{\\tau}$, i.e. $\\rho \\propto T^4$. Geodesic motion of photons within the holostar preserves the Planck-distribution. The radial position of an observer can be determined via the local mass-density, the local radiation-temperature, the local entropy-density or the local Hubble-flow. Using the experimentally determined values for the matter-density of the universe, the Hubble-value and the CMBR-temperature, values of $r$ between $8.06$ and $9.18 \\cdot 10^{60} r_{Pl}$ are calculated, i.e values very close to the current radius and age of the universe. Therefore the holographic solution might serve as an alternative model for a universe with an overall negative (string type) equation of state, without need for a cosmological constant. The holostar as a model for a black hole or the universe contains no free parameters: The holostar metric and fields, as well as the initial conditions for geodesic motion are completely fixed and arise naturally from the solution. An analysis of the characteristic properties of geodesic motion in the interior space-time points to the possibility, that the holostar-solution might contribute substantially to the understanding of the phenomena in our universe. The holostar model of the universe is free of most of the problems of the standard cosmological models, such as the \"cosmic-coincidence-problem\", the \"flatness-problem\", the \"horizon-problem\", the \"cosmological-constant problem\" etc. . The cosmological constant is exactly zero. There is no horizon-problem, as the co-moving distance $r$ evolves exactly proportional to the Hubble-radius ($r \\propto \\tau$ for $r \\gg r_{Pl}$). Inflation is not needed. There is no initial singularity. The expansion (= radially outward directed geodesic motion) starts out from a Plank-size region at the Planck-temperature, which contains at most one \"particle\" with roughly 1/8 to 1/5 of the Planck-mass. The maximum angular separation of particles starting out from the center is limited to roughly $60^\\circ$, which could explain the low quadrupole and octupole-moments in the CMBR-power spectrum. The relation $r \\propto \\tau$ for $r \\gg r_{Pl}$, whose fundamental origin can be traced to the zero active gravitational mass-density of the string-type matter in the interior space-time, can be interpreted in terms of a permanently unaccelerated expansion, from which $H \\tau = 1$ follows. This genuine prediction of the holostar model is very close to the observations, which give values between $0.98$ and $1.04$. Permanently undecelerated expansion is also compatible with the luminosity-redshift relation derived from the most recent supernova-measurements, although the experimental results favor the concordance $\\Lambda$CDM-model over the holostar-model at roughly $1 \\sigma$ confidence level. The Hubble value in the holostar-model of the universe turns out lower than in the concordance model. $H \\simeq 1/r = 62.35 \\, (km/s) / MPc$ is predicted. This puts $H$ into the range of the other absolute measurements, which consistently give values $H \\approx 60 \\, (km/s) / MPc$ and is compatible with the recent supernova-data, which favor values in the range $H \\approx 60-65$. Geodesic motion of particles in the holostar space-time preserves the relative energy-densities of the different particle species (not their number-densities!), from which a baryon-to photon ratio of roughly $\\eta \\approx 10^{-9}$ at $T_{CMBR} = 2.725 \\, K \\approx 10^{-9} (m_e c^2 / k)$ is predicted. The holographic solution also admits microscopic self-gravitating objects with a surface area of roughly the Planck-area and zero gravitating mass. ", "introduction": "Introduction:} In a series of recent papers new interest has grown in the problem of finding the most general solution to the spherically symmetric equations of general relativity, including matter. Many of these papers deal with anisotropic matter states.\\footnote{Relevant contributions in the recent past (most likely not a complete list of relevant references) can be found in the papers of \\cite{Barve/Witten, Burinskii, Dev/00, Dev/03, Dymnikova, Elizalde, Gair, Goncalves, Giambo, Hernandez, Herrera, Herrera/2002, Ivanov, Mak, McManus/Coley, Mielke/BosonStars, petri/bh, Salgado, Singh/Witten} and the references given therein.} Anisotropic matter - in a spherically symmetric context - is a (new) state of matter, for which the principal pressure components in the radial and tangential directions differ. Note that an anisotropic pressure is fully compatible with spherical symmetry, a fact that appears to have been overlooked by some of the old papers. One of the causes for this newly awakened interest could be the realization, that models with anisotropic pressure appear to have the potential to soften up the conditions under which spherically symmetric collapse necessarily proceeds to a singularity.\\footnote{See for example \\cite{Singh/Witten}, who noted that under certain conditions a finite region near the center necessarily expands outward, if collapse begins from rest.} Another motivation for the renewed interest might be the prospect of the new physics that will have to be developed in order to understand the peculiar properties of matter in a state of highly anisotropic pressure and to determine the conditions according to which such matter-states develop. The study of anisotropic matter states on a large scale might also become relevant from recent cosmological observations. It is well known, that the cosmic microwave background radiation (CMBR) contains a dipole with a direction pointing roughly to the Virgo cluster. The common explanation for this anisotropy is the relative motion of the earth with respect to the preferred rest-frame of the universe, which is identified with the frame in which the CMBR appears spherically symmetric. Yet despite years of research the physical cause for this peculiar motion has remained unclear. The observations have failed to deliver independent conclusive evidence for a sufficiently large mass concentration in the direction of Coma/Virgo. In a recent paper the question was raised, whether the universe might exhibit an intrinsic anisotropy at large scales \\cite{Ralston}. Not only the CMBR-dipole, but also the quadrupole and the octupole coefficients of the CMBR-multipole expansion single out preferred axes, which all point in the same direction as the dipole, i.e. roughly to the Virgo direction.\\footnote{In order to get directional information from the higher multipoles, the authors analyzed not only the (directionless) $l$-terms of the multipole expansion of the CBMR-powerspectrum in spherical harmonics, but also the $m$-terms.} In an earlier paper the same authors found an alignment of optical and radio polarization data with respect to Virgo \\cite{Jain/polarization, Jain}.\\footnote{In \\cite{Jain/polarization} the authors analyzed the statistics of offset angles of radio galaxy symmetry axes relative to their average polarization angles, indicating an anisotropy for the propagation of radiation on cosmological scales, which lies in the direction of Virgo. Another preferred axis, parallel to the CMBR-dipole within the measurement errors, was identified in \\cite{Jain}. Here the optical polarization data from cosmologically distant and widely separated quasars showed an improbable degree of coherence. A significant clustering of polarization coherence in large patches in the sky was identified, with the axis of correlation again lying in the direction of Virgo.} In \\cite{Ralston} a Bayesian analysis was performed with the result, that the common alignment of five different axes appears very unlikely in the context of the Standard Cosmological Model of a homogeneously expanding Friedman Robertson-Walker universe. The authors argue, that within the standard big bang hypothesis one rather expects an isotropic distribution of the different multipole alignment axes for the post-inflationary state. As the experimental evidence rather points to the contrary the authors conclude, that we might be living in a universe that exhibits an intrinsic anisotropic on very large scales. In a very recent paper another team of authors pointed out, that the quadrupole and (three) octupole planes are correlated with 99.97 \\% confidence level and that the alignment of the normals of these planes with the cosmological dipole and the equinoxes is inconsistent with Gaussian random skies at 99.8 \\% confidence level \\cite{Schwarz2004}. Anisotropic matter states are also predicted by string theory. The equation of state for a cosmic string or for a 2D-membrane is naturally and necessarily anisotropic. The interior matter-state of the solution discussed in this paper has a definite string interpretation: It is that of a collection of radially outlayed strings, attached to a spherical 2D-boundary membrane. The string tension falls off as $1/r^2$. The strings are \"densely\" packed in the sense, that the transverse extension of the strings amounts to exactly one Planck area \\cite{petri/string}. The absence of singularities and trapped surfaces in the holostar space-time is compatible with a very recent result in string theory. According to \\cite{Senovilla} the so called VSI space-times\\footnote{VSI = vanishing scalar invariants. VSI-space-times are solutions where all scalar invariants vanish} - exact solutions of string theory - are incompatible with trapped surfaces. The proof is quite general. It is based on geometric arguments and doesn't require the field equations. If this result is confirmed, string theory might turn out to be incompatible with trapped surfaces and - consequentially - classical space-time singularities, such as can be found in the classical black hole solutions.\\footnote{In \\cite{Chinea2004} it was shown, that trapped surfaces (or more generally: the occurrence of any causal trapped set in the space-time) are an {\\em essential} requirement for the singularity theorems: Neither the energy-conditions nor the causality conditions alone or in combination lead to a singularity. Whereas the positivity of the energy and the principle of causality are fundamental requirements for a self-consistent physical theory, trapped surfaces are not necessary for a self-consistent physical description. The {\\em assumption} that a {\\em physically realistic} space-time should develop a trapped surface (or a trapped causal set) at some particular time is the most questionable of the assumptions underlying the singularity theorems. One can regard it as the key assumption on which the common belief is based, that a physically realistic space-time must contain singularities. But this assumption has not been proven. It was formulated when the only solutions of physical relevance that were known at that time {\\em all} contained trapped surfaces: the black hole solutions. Now we know of physically relevant solutions {\\em without} trapped surfaces and singularities. The claim, that a physically realistic space-time must necessarily contain trapped surfaces (and therefore singularities), must be viewed in the proper historical context. If history had taken a different route, chances are, that today's claims would have been quite different: If realistic singularity free solutions to the field equations had been known at the time the singularity theorems were formulated, one would rather have argued, that singularities are unphysical and therefore trapped surfaces cannot be elements of a physically realistic space-time. In view of the new singularity free solutions it is appropriate to exercise some caution in raising an assumption to the status of unquestionable physical truth, as can be found occasionally in today's scientific literature. If we are honest, we must admit that despite decades of research we still don't know, what properties a \"physically realistic\" space-time should actually have. In fact, our preconceptions about how a physically realistic space-time should \"look\" like have changed dramatically in the course of history. A most radical change occurred in the recent past: The luminosity dependence of distant supernovae on red-shift makes it quite clear, that we are not living in a dust-type universe with equation of state $P \\approx 0$, as had been thought decades before, but rather in a universe which consists mostly out of (cold) \"dark matter\" and \"dark energy\". Although we neither know what the \"dark matter\" is, and even less what the \"dark energy\" could be, the equation of state for the large scale distribution of mass-energy in the universe has been measured fairly accurately: It is of the form $\\rho +P \\approx 0$, which is compatible with vacuum-type, string-type or domainwall-type matter.} Therefore the result reported in \\cite{Senovilla} has two possible outcomes: Either string theory is the - essentially - correct description for a physically realistic space-time in the high {\\em and} the low energy limit. Then trapped surfaces, the classical vacuum black hole solutions and - most likely - space-time singularities (of \"macroscopic scale\") cannot be part of a physically realistic description of the world we live in. We will have to search for other solutions of the classical field equations without singularities and trapped surfaces, preferentially with a strong string character. The holostar is the simplest member of such a class of solutions. The other outcome - which might be the preferred scenario by the 1000 or more researchers who invested much time and effort in the study of singularities, event horizons, apparent horizons and trapped surfaces - might be, that despite their highly undesirable properties event horizons, trapped surfaces and singularities are real. Then string theory cannot fulfill its main objective to provide a unified description for all forces, including gravity, on all energy scales. We will have to look for another approach. Whatever the final outcome is going to be, for the present time we have to work with the theories and solutions we know. In this paper the geometric properties of the so called holographic solution, an exact solution of the Einstein field equations, are studied in somewhat greater detail. The holographic solution turns out to be a particularly simple model for a spherically symmetric, self gravitating system with a highly - in fact maximally - anisotropic, string-type pressure. Despite its simplicity and its lack of adjustable parameters, the holostar appears very well suited to explain many of the phenomena encountered in various self gravitating systems, such as black holes, the universe and - possibly - even elementary particles. The paper is divided into three sections. In the following principal section some characteristic properties of the holographic solution are derived. In the next section these properties are compared to the properties of the most fundamental objects of nature that are known so far, i.e. elementary particles, black holes and the universe. The question, whether the holostar can serve as an alternative, unified model for these fundamental objects (within the limitations of a classical description!) will be discussed. The paper closes with a discussion and outlook. ", "conclusions": "The holostar solution appears as an attractive alternative for a black hole. It has a surface temperature, which - measured at spatial infinity - is proportional to the Hawking temperature. Its total number of interior relativistic particles is proportional to its proper surface area, which can be interpreted as evidence that the Hawking entropy is of microscopic-statistical origin and the holographic principle is valid in classical general relativity for self gravitating objects of arbitrary size. In contrast to a black hole, the holostar has no event horizon, no trapped surfaces and no central singularity, so there is no information paradox and no breakdown of predictability. The interior matter-state of a holostar consists of string-type matter, which extends to the holostar's boundary membrane, situated roughly 2 Planck coordinate lengths outside the holostar's gravitational radius. The string tension at the boundary is given by $\\mu = 1 / (8 \\pi r^2)$, which is inverse proportional to the mean string length $R = r^2 / ( 2 r_0)$. This result agrees with a very recent result in string theory \\cite{Mathur/2004}. Furthermore, the holostar solution has some potential to serve as an alternative model for a universe with anisotropic negative pressure, without need for a cosmological constant. It also admits microscopic self-gravitating objects with a surface area of roughly the Planck-area and zero gravitating mass. The remarkable properties of the holostar solution and its high degree of self-consistency should make it an object of considerable interest for future research. A field of research which presents itself immediately is the the generalization of the holostar solution to the rotating and / or charged case. The derivation of the charged holostar solution is straight forward and is discussed in \\cite{petri/charge}. A generalization to the case of a rotating body, including spin (and charge), will be an interesting topic of future research. An accurate description of the thermodynamic properties of the holostar solution is of considerable interest. In \\cite{petri/thermo} the entropy/area law and the temperature-law for the interior holostar matter state are put on a sound foundation. The thermodynamic analysis allows one to relate the Hawking temperature (at spatial infinity) to the local temperature and energy-density in the interior holostar space-time. Using this relation the Hawking prediction is verified to an accuracy of roughly 1\\%. Another valuable field of research will be the examination, whether the holostar solution can serve as an alternative model for the universe. The holostar solution appears to have the potential to solve many problems of the standard cosmological models, such as the horizon-problem, the cosmological constant problem, the problem of structure formation from the small density perturbations in the CMBR, etc. . It provides an explanation for the numerical value of the baryon to photon ratio $\\eta$. Even some of the more recent observational results, such as the missing angular correlation of the CMBR-fluctuations at angular separations larger than $60^\\circ$, are explainable in the context of the holostar space-time. On the other hand, it is far from clear whether the holostar solution will ever be able to explain the observed abundances of the light elements, especially Deuterium and Lithium, in a convincing fashion, such as the Standard Cosmological Model can. Nucleosynthesis in a holostar universe will be a demanding challenge. If the holostar can pass this test, it should open up a new field of interesting research in cosmology and particle physics. Quite likely chemical potentials and supersymmetry will play an important role in the holostar universe, at least at temperatures at the GUT-scale. Lastly, the properties of the membrane, how it is formed, how it is composed and how it manages to maintain its two-dimensional structure will be an interesting research topic, if the holostar turns out to be a realistic alternative for a black hole. The study of anisotropic string type matter states in high gravitational fields should provide fruitful as well." }, "0405/astro-ph0405054_arXiv.txt": { "abstract": "{The star HD~219542~B has been reported by us (Desidera et al. 2003) to show low-amplitude radial velocity variations that could be due to the presence of a Saturn-mass planetary companion or to stellar activity phenomena. In this letter we present the results of the continuation of the radial velocity monitoring as well as a discussion of literature determinations of the chromospheric activity of the star (Wright et al.~2004). These new data indicate that the observed radial velocity variations are likely related to stellar activity. In particular, there are indications that HD~219542~B underwent a phase of enhanced stellar activity in 2002 while the activity level has been lower in both 2001 and 2003. Our 2003 radial velocity measurements now deviate from our preliminary orbital solution and the peak in the power spectrum at the proposed planet period is severely reduced by the inclusion of the new data. We therefore dismiss the planet hypothesis as the cause of the radial velocity variations. ", "introduction": "\\label{s:intro} In Desidera et al.~(\\cite{hd219542}, hereafter Paper I) we presented high precision radial velocity (hereafter RV) monitoring of the components of the wide binary system HD~219542. This pair is part of the sample of wide binaries currently under monitoring at TNG using the high resolution spectrograph SARG (Gratton et al.~\\cite{papersarg}). We have found evidence for the presence of low amplitude RV variations on the secondary HD~219542~B with a period of 112 days at a confidence level of 96-97\\%. These RV variations could be due to a Saturn-mass planet orbiting at 0.46~AU or to stellar activity. The relatively low statistical confidence of the proposed planetary orbit as well as the possible presence of stellar activity indicate that a confirmation is required before we can classify of HD~219542~B as a bona fide planet host star. It is well known that stellar activity induces distortions of the profile of the spectral lines that could be seen as RV variations and then mimic the occurrence of companions orbiting the target (see e.g. Saar et al.~\\cite{saar}). These spurious RV variations may have amplitudes and timescales comparable to those induced by giant planets, making challenging the search for planets around active stars. The controversial case of HD~192263 (Henry et al.~\\cite{henry02}; Santos et al.~\\cite{santos_hd192263}) is worth of mention in this context. Two different components of the magnetic activity phenomena are of concern for planet searches. Star spots alter the profile of spectral lines, as well known from Doppler imaging studies of rapidly rotating spotted stars (see e.g. Rice \\cite{rice}; Strassmeier \\cite{strassmeier}). For slowly rotating stars the distortions of line profile are more subtle but nevertheless sufficient to be detected as spurious RV variations (Hatzes \\cite{hatzes02}). The RV variations resulting from the presence of surface features typically follow the time scales of the rotational period of the star (a few days for the active stars for which such signal is more easily detectable), but for the long term and sparse sampling typical of planet search surveys no clear periodicities are often present, because of the limited lifetime of such features. RV variations caused by star spots are usually correlated to photometric variations (Paulson et al.~\\cite{paulson04}; Queloz et al.~\\cite{hd166435}). The second contribution is represented by plages, that cause a change of the shape of spectral lines, mostly because of the alteration of the granulation pattern (Saar \\cite{saar03}; K\\\"urster et al.~\\cite{barnard}). The variations of the area covered by plages along the magnetic cycle and/or the rotational period then cause RVs variations, correlated with chromospheric emission. For some low activity stars, for which the effects of rotational modulations are lower than those of the long term magnetic cycle, a fairly good correlation between RVs and chromospheric emission can be found (Saar \\& Fischer \\cite{saarfischer}). In order to disentangle the origin of RV variations (keplerian orbital motion {\\it vs} activity jitter) basically three approaches can be pursued. \\begin{enumerate} \\item To directly search for the presence of distortions of line profiles (ideally on the same spectra on which RVs are derived), as done by e.g. Hatzes et al.~(\\cite{hatzes98}) for 51 Peg and Queloz et al.~(\\cite{hd166435}) for HD~166435 \\item To obtain measurements of stellar activity (chromospheric emission, photometry), possibly simultaneous to RVs, to search for correlation between RVs and activity. Note that one single activity diagnostics may not be enough since photometry and chromospheric emission are mostly sensitive to different components of magnetic activity. \\item To continue the RV monitoring of the object. In fact, stellar activity typically is not stable on yearly timescales, so that an activity signal is not expected to maintain the same phase and amplitude with time (see Queloz et al.~\\cite{hd166435}) \\end{enumerate} For our candidate we followed the second and third approaches, collecting new RV data and considering literature determinations of the chromospheric emission of the star (Wright et al.~\\cite{rhk_keck}), published after Paper I, to study the evolution of its activity level. The line bisector variations corresponding to the observed RV variations are below our detectability threshold, as discussed in Paper I. ", "conclusions": "The continuation of the radial velocity monitoring and the multi-epoch measurements of the Ca II H\\&K emission indicate that the low amplitude RV variations of HD~219542~B presented in Paper I are likely due to stellar activity. This star should therefore be removed from the list of extrasolar planet host stars. The available data suggest that the star underwent a relatively short-lived phase of enhanced activity during the 2002 season. This study confirms the relevance of the activity-related phenomena in the RV planet searches and, on the other hand, the great impact of the high precision Doppler surveys in improving our understanding of the stellar activity cycles. Even for stars with modest activity level, the discovery of planets with amplitude of 10-15 m/s requires a long term monitoring to check for the stability of the signal as well as ancillary measurements of activity indicators (see e.g. Hatzes et al.~\\cite{epseri})." }, "0405/astro-ph0405112_arXiv.txt": { "abstract": "{A study of the gas content in 1038 interacting galaxies, essentially selected from Arp, Arp and Madore, Vorontsov-Velyaminov catalogues and some of the published literature, is presented here. The data on the interstellar medium have been extracted from a number of sources in the literature and compared with a sample of 1916 normal galaxies. The mean values for each of the different ISM tracers (FIR, 21 cm, CO lines, X-ray) have been estimated by means of survival analysis techniques, in order to take into account the presence of upper limits. From the data it appears that interacting galaxies have a higher gas content than normal ones. Galaxies classified as ellipticals have both a dust and gas content one order of magnitude higher than normal. Spirals have in most part a normal dust and HI content but an higher molecular gas mass. The X-ray luminosity also appears higher than that of normal galaxies of same morphological type, both including or excluding AGNs. We considered the alternative possibilities that the molecular gas excess may derive from the existence of tidal torques which produce gas infall from the surrounding regions or from a different metallicity which affects the X conversion factor between the observed CO line luminosity and the H$_2$ calculated mass. According to our tests, it appears that interacting galaxies possess a higher molecular mass than normal galaxies but with a similar star formation efficiency. \\keywords {catalogues -- galaxies: ISM -- galaxies: Interacting}} ", "introduction": "It is known that interactions between galaxies and environment play an important role in determining the internal galaxy structure. Gas accretion may produces central black holes \\citep{bh}, polar rings \\citep{pr} or counter-rotations \\citep{crgg, crbis}. Gas stripping in clusters may deplete a spiral galaxy transforming it into an S0 \\citep{dressler97}, while massive collisions may destroy the whole galaxy structure creating giant triaxial ellipticals \\citep{bendo}. According to numerical models, bars and rings may also be generated by a close encounter between galaxies \\citep{lia}. Interaction also affects other galaxy properties, such as the star formation rate \\citep{sanders,sage88,combes}, visible as a strong increase of infrared luminosity. Inside galaxies, gas is expected to reflect the effects of the interaction more strongly than stars. A recent study on the interstellar medium (ISM) in 104 peculiar galaxies \\citep{polar} shows that polar ring galaxies have a gas content one order of magnitude higher than normal ones. In many cases, the higher gas content is visible even if the interaction with the environment has ceased for long time. We wonder if a peculiar gas content is detectable in those galaxies where interaction is currently ongoing, such as some systems contained in the \\citet{vv}, \\citet{arp} and \\citet{am} catalogues. Here we present the results of a selection of 1038 objects extracted from these catalogues. For these galaxies data on dust, HI, molecular gas and X-ray luminosities are available in the literature. The properties of the different components of the interstellar medium are studied here and compared with those of a sample of 1916 normal galaxies \\citep{normal}. \\setcounter{table}{1} \\begin{table} \\caption{Number of interacting galaxies observed (N) and detected (N$_d$) according to the different ISM tracers. The column labeled `All' represents the number of galaxies present in each morphological type bin. The type 11 is assigned here to galaxies whose structure is not attributed to a specific morphological type in LEDA. } \\tabcolsep 0.15truecm \\begin{center} \\begin{tabular}{lrrrrrr} \\hline \\hline &\t\t& All\t & \\multicolumn{1}{c}{Dust} &\t\\multicolumn{1}{c}{HI}\t& \\multicolumn{1}{c}{Mol} &\t\\multicolumn{1}{c}{X-ray}\t\\\\ Type\t&\tt\t& \t & N/N$_d$\t& N/N$_d$\t& N/N$_d$\t& N/N$_d$\t\\\\ \\hline E \t&\t-5 \t& 22 & 15/8 \t& \t12/7 \t& \t3/0 \t& \t13/11\t\\\\ E$^+$ & \t-4 \t& 21 & 18/12 \t& \t6/3 \t& \t3/1 \t&\t11/10\t\\\\ E/S0 & \t-3 \t& 18 & 13/9 \t& \t9/5 \t& \t4/1 \t&\t9/4 \t\\\\ S0 & \t\t-2 \t& 47 & 36/31 \t& \t21/14\t& \t8/2 \t& \t11/7 \t\\\\ S0$^+$ &\t-1 \t& 40 & 35/31 \t& \t18/15\t& \t3/2 \t& \t4/1 \t\\\\ S0/a & \t 0 \t& 36 & 31/25 \t& \t14/14\t& \t7/4 \t& \t3/1 \t\\\\ Sa \t& \t 1 \t& 82 & 75/71 \t& \t44/41\t& \t12/12\t& \t9/5 \t\\\\ Sab \t&\t 2 \t& 78 & 67/63 \t& \t42/42\t& \t10/9 \t& \t6/3 \t\\\\ Sb \t& \t 3 \t& 142 & 128/127 \t& \t75/72\t& \t29/23\t& \t8/5 \t\\\\ Sbc \t& \t4 \t& 177 & 169/164 \t& \t68/64\t& \t16/15\t& \t10/8 \t\\\\ Sc+S? & \t5 \t& 148 & 129/124 \t& \t89/87\t& \t19/17\t& \t10/3 \t\\\\ Sc \t& \t6 \t& 74 & 63/58 \t& \t50/48\t& \t12/11\t& \t8/2 \t\\\\ Scd \t& \t7 \t& 33 & 31/30 \t& \t20/20\t& \t4/4 \t& \t5/4 \t\\\\ Sd \t& \t8 \t& 30 & 23/20 \t& \t26/26\t& \t4/2 \t& \t2/1 \t\\\\ Sm \t& \t9 \t& 30 & 22/19 \t& \t25/24\t& \t10/6 \t& \t4/3 \t\\\\ Irr \t& \t10 \t& 35 & 23/23 \t& \t30/29\t& \t5/5 \t& \t4/3 \t\\\\ ? \t& \t11 \t& 25 & 22/21 \t& \t7/7 \t& \t4/4 \t& \t1/0 \\\\ \\hline Total & \t\t& 1038 & 900/836 \t& 556/518 \t& 153/118 \t& \t118/71 \\\\ \\hline \\hline \\end{tabular} \\end{center} \\label{numbers} \\end{table} ", "conclusions": "The category of galaxies considered in this paper is a zoo containing several types of objects, going from galaxies observed during a close encounter (e.g. M51) to merging systems (e.g The Antennae) or galaxies relatively isolated but still perturbed after a past encounter (e.g. M82). The spread of values found here and in the literature may derives from the fact that different kind of perturbations generate different effects in galaxies. We find that the morphological type also have a role in determining the properties of these stellar systems and that may be not correct to mix together interacting galaxies of all morphological types. In addition, the comparison with the literature is complicated by the fact that very often authors use as `normal' galaxies a list of objects that contains known cases of peculiar galaxies. For instance, in Table 3 of \\citet{devereux97} is visible NGC 660, a disrupted polar ring galaxy. \\citet{HG} include several Arp and VV galaxies while \\citet{casoli} present a sample containing many dynamically peculiar galaxies. The data of these papers have been used by other authors to deduce properties of peculiar galaxies in comparison with `normal' ones. The sample of normal galaxies by \\citet{normal} used here may also contains galaxies that in the future may be discovered as 'peculiar', but have been already cleaned by all the known cases of interacting and dynamically peculiar galaxies known. From our data, we see that interacting or disturbed galaxies have always more gas than normal ones, and that this gas is mainly in the form of atomic hydrogen in early type systems, while it is in molecular form for galaxies of latest types. All have a higher X-ray luminosity. Because of these differences, we shall discuss the early type systems separately from the spiral ones, comparing our finding with that from literature. A galaxy that appears now as an elliptical may be in the final stage of a merging process, after the complete fusion of the stellar content of two galaxies. Alternatively, it may be a galaxy deprived of the gas because of a close encounter with a more massive one. In both processes, one may expect that the pre-existing gas in the two systems would be heated, with the conversion of molecular gas in atomic form and by the creation of an X-ray component. This scenario agrees with the higher HI and X-ray content found in interacting early type galaxies. A more complicated scenario appears from the data concerning galaxies classified as spirals. In these systems the mean HI and dust content seems to be similar to that of an unperturbed, normal galaxy of the same morphological type. In addition to the excess in X-ray luminosity, these galaxies seem to have a molecular gas content always one order of magnitude higher that that of normal galaxies. There are two possible explanation for this effect. The first possibility is that during the collision the galaxies stimulate a gas inflow by means of gravitational torques, that enhance the CO luminosity because of a massive accumulation of molecular gas. This hypothesis has been presented by \\citet{combes} to explain the observed characteristics of a sample of 51 paired galaxies, whose interacting subsample appears more luminous in CO line than the remaining binaries with higher separation. We tested the idea that the gas is simply exchanged by the two components of the same pair, enriching one of them and depauperating the other. We call these galaxies primary and secondary respectively. In our sample there are 179 galaxies in interacting pairs or triplets where the single galaxies are separated and have literature data at 21cm or CO lines. Plotting in Figure \\ref{coppie} (upper panel) the values of log \\mhi/\\LB\\ of these galaxies only and identifying with different symbols the primary and the secondary member, one can see that spirals which are primary members have -on mean- higher values than normal galaxies of same morphological type. On the contrary, the secondary members, of whatever morphological type, have values generally lower, below that of normal galaxies. This is in agreement with the idea that part of atomic gas has been transferred between them. However, this is not true when log \\mmol/\\LB\\ values are considered (Figure \\ref{coppie}, lower panel): in this case, the primary members of the pair show an excess of molecular mass and the secondary members have in general values near or over the mean curve of normal galaxies. We may suppose that less dense, atomic gas, may be exchanged during an encounter more easily than the molecular gas, condensed in clouds. The existence of a molecular mass excess in these paired galaxies is then difficult to explain supposing a simple transfer during the close encounter. \\begin{figure} \\resizebox{9cm}{!}{\\includegraphics[angle=0]{0283fig7.eps}} \\caption{Distribution of the representative points for 179 galaxies in pairs or triplets present in our sample for which detections or upper limits are available. The higher mass member (primary) is shown with a full diamond while the lower mass ones (secondary) are indicated with a cross. } \\label{coppie} \\end{figure} The second, completely different possibility is that the molecular gas excess is not real but due to different physical conditions inside the galaxy. In the totality of cases presented here the \\mmol\\ is determined from CO line luminosities, transformed in solar masses adopting a X conversion factor, as explained in Section \\ref{methods}. This value depends on the metallicities or on the physical conditions of the molecular gas and in different kind of galaxies may assume a wide range of values (see the discussion in Section \\ref{sfr}). An overestimate of the molecular mass may derives by the assumption of a X factor lower than real. among the various possibility, the presence of gas of high temperature may overestimate the H$_2$ mass. However, in some well known cases present in our sample, such as the already cited `Antennae' or `Atoms for Peace' (PGC 68612), believed to be the result of a two spirals merger, the H$_2$ mass excess is accompanied by an HI mass excess. Similarly the `Whirlpool galaxy' M51 and its satellite (PGC 47404 and PGC 47413), show an excess of gas mass. In cases like these, we may think that the gas excess is real for the molecular component like it is for the atomic one. Another suggestion on the reality of the higher molecular masses comes from the analysis of the star formation rates and efficiencies: interacting galaxies appear more luminous in the infrared (Fig. \\ref{LFIR}) because of the higher number of newly formed stars heating the dust \\citep{tronson}. This higher star formation rate does not correspond to a different efficiency in star formation per unit of calculated H$_2$ mass. This may be explained if the higher star formation is simply due to a higher molecular gas quantity. Higher H$_2$ mass and concentration in perturbed galaxies was already suggested by \\citet{braine}. In our data, the large spread of data in Figure \\ref{SFE} may indicate that the X conversion factor is different between galaxies but even in this hypothesis, the range of variation inside interacting galaxies does not appear different from that of normal, non interacting stellar systems. Our conclusion are different with the higher SFE found in interacting galaxies by previous authors \\citep{sage88,combes,horellou3}. As discussed before, their comparison sample contains peculiar galaxies and the large spread of values found in our sample cover the differences found by them. We must conclude that, with the present data, our sample of interacting galaxies appears richer in molecular gas and in some cases also in atomic gas. It is not clear from where this mass excess originates. An excess of gas, both atomic and molecular, of one order of magnitude is known in the literature for another class of peculiar objects: the polar ring galaxies \\citep{polar}. In that case, the matter forming the ring is accreted from outside. May this also be the case of interacting galaxies?" }, "0405/astro-ph0405438_arXiv.txt": { "abstract": "{Brown dwarfs}\\index{brown dwarfs} bridge the gap between the stellar and planetary mass regimes. Evolving from conditions very similar to the lowest-mass stars, the atmospheres of older brown dwarfs closely resemble those expected in close-in {extrasolar giant planets}\\index{extrasolar giant planets}, and with cooler BDs still being discovered, more and more approach the properties of gas giants at wider separation. Interpreting the spectra of BDs is therefore a crucial step towards understanding and predicting the spectral and thermal properties of EGPs. Essential properties of substellar atmospheres are massive molecular line-blanketing and the condensation of species with decreasing {$T_{\\mathrm{eff}}$}, changing the chemical equilibrium composition and causing absorption from dust grains. More complex details involve the distribution of dust clouds over the surface giving rise to temporal variability, and possible deviations from chemical equilibrium conditions. In the case of close-in EGPs and some BDs in binary systems, the effect of irradiation from the primary significantly affects the spectral properties and thermal evolution. ", "introduction": "Observational efforts during the past decade have brought marked progress in characterising the lower end of the Main Sequence down to and beyond the hydrogen-burning limit, and identifying a class of unambiguously {substellar}\\index{substellar} objects. Brown dwarfs (BDs) are commonly defined as compact objects with a mass below the minimum for sustaining equilibrium hydrogen fusion ($\\sim\\!0.07 M_\\odot$). Since the detection of the first bona-fide BDs \\cite{reboloPleiade,nakaGl229B,broGl229B} and the discovery of a Jupiter-mass companion to 51\\,Peg by Mayor \\& Queloz \\cite{mayor51Peg}, the field of substellar astronomy has seen the direct detection of more than 300 ultracool dwarfs and the indirect detection of 110 planetary mass objects in orbit around other stars. While there is no agreement yet whether brown dwarfs should be distinguished from extrasolar giant planets (EGPs) based on their formation history, or by defining the minimum mass for {deuterium burning}\\index{deuterium burning} ($\\sim\\!13 M_\\mathrm{J}$) as the lower limit of the BD mass range, in terms of atmospheric properties and spectroscopic appearance there is a smooth transition, and some overlap, between both classes of objects, regardless of definition. The spectroscopic characterisation of these sources % of 2200\\,K\\,$>T_\\mathrm{eff}>$\\,700\\,K has prompted the introduction of the two new {spectral classes L and T}. While the original description of the L sequence was based on optical spectral properties \\cite{mbdfLdwarf97,krlLdwarf99}, the definition of a classification scheme for the latest L-type\\index{L dwarfs} and the {T dwarfs}\\index{T dwarfs} has required the use of near-IR spectra, owing to the fact that the lion's share of emitted flux and of spectroscopic characteristics in these objects is found at $\\lambda>0.95$\\,\\umu m \\cite{geb02,adam02a}. While the youngest / most massive brown dwarfs jointly populate the early L (and late M) classes together with the least massive hydrogen-burning (very low mass-, or VLM-) stars, more evolved brown dwarfs of late-L type have effective temperatures similar to those of the EGPs closest-in to their primary star (\\lq \\lq {Hot Jupiters}\\index{Hot Jupiters}\", e.\\,g.\\ 51~Peg, HD~209458 \\cite{evolPlanets}). The coolest T dwarfs detected to date \\cite{tomgGl570,2MASS04pap} exhibit atmospheric properties close to those anticipated in EGPs at larger separation, and thus are crucial test cases for our understanding of planetary atmospheres. The spectral appearance of substellar objects, proceeding from the lowest-mass main-sequence stars to lower effective temperatures is mainly characterised by the onset of condensate formation in the latest {M dwarfs}\\index{M dwarfs}, leading to strongly dust-dominated atmospheres of early L-types with very red $J\\!H\\!K$-colours. From mid-L to early T-types this trend is reversed with a decreasing signature of dust, strengthening of water vapour absorption and the appearance of methane. The strongest CH$_4$ bands at 3.3 and 2.2\\,{\\umu m} have been observed in dwarfs as early as L5 \\cite{2000ApJ...541L..75N}. The L/T transition is thus now defined by the first appearance of the weaker band at 1.6\\,{\\umu m} \\cite{geb02,adam02a}. While the near-IR spectra of T dwarfs are turning blue again with decreasing $T_\\mathrm{eff}$, caused by the deepening molecular bands and {collision-induced absorption}\\index{collision-induced absorption} (CIA) of H$_2$ in the $K$-band, their optical-infrared colours become extremely red due to the massively pressure-broadened resonance lines of Na\\,I and K\\,I. Chemical interaction, the complex spectra of molecules, the physics of {dust formation}\\index{dust formation} and the treatment of {line broadening}\\index{line broadening} in high density conditions make spectral modelling of brown dwarfs and giant planets a challenging task. With advancements in computational techniques and availability of better chemical and physical input data over the years, models of the atmospheres and interiors, \\cite{1993ApJ...406..158B,faphh95,tsujiDust96,evolDust} have been able to keep up with observational progress in their ability to describe the global properties of brown dwarfs. In this review, we present an introduction into dwarf atmosphere modelling with the \\texttt{PHOENIX} code in Sect.~\\ref{phoenix}, followed by an overview of the current status of models and recent developments in Sect.~\\ref{status}. We discuss the special case of modelling the {irradiated atmospheres}\\index{irradiated atmospheres} of close-in extrasolar giant planets in Sect.~\\ref{planets} and discuss the outlook to possible direct observations of EGPs. \\markboth{D.~Homeier et al.}{Brown Dwarfs and Hot Jupiters} ", "conclusions": "Ten years of observational and theoretical work have brought us thorough understanding of the basic physical properties of brown dwarfs. While many details, especially regarding non-equilibrium processes, and temporal and spatial variability of the atmosphere, still need to be addressed, current models provide a solid foundation for understanding extrasolar giant planets, and predicting their observable signatures. BD and EGP (unirradiated as well as \\lq \\lq Hot Jupiter''-like) model atmospheres, thermal profiles, synthetic spectra, photometry, and evolution calculations are available for all evolutionary stages from a few Myrs from our web site at: \\href{http://perso.ens-lyon.fr/france.allard/}{http://perso.ens-lyon.fr/france.allard/}. \\subsection*" }, "0405/astro-ph0405324_arXiv.txt": { "abstract": "We present intermediate resolution near-IR long-slit spectroscopic data for the nearby radio galaxy Cygnus A (3C 405) (obtained with the NIRSPEC spectrograph on the Keck II telescope). The data reveal considerable complexity in the near-IR emission line kinematics, including line splittings of 200-350 km s$^{-1}$ and a mixture of narrow (FWHM $\\sim$200 km s$^{-1}$) and broad (FWHM $\\sim$700 km s$^{-1}$) components to the emission lines. It is notable that the Pa$\\alpha$ and H$_2$ emission lines show markedly different kinematics, both on- and off-nucleus. Overall, the data provide evidence for the presence of a giant molecular cloud falling through the heart of the Cygnus A host galaxy, the motion of which is not driven by the AGN itself. We suggest that this cloud may be connected to the triggering of the activity in this highly powerful AGN. We also detect split H$_2$ components on the nucleus that are likely to originate in the circum-nuclear torus. ", "introduction": "The prototype radio galaxy Cygnus A (3C 405) is the most luminous extragalactic radio source in the local universe and the only nearby (z=0.0558) radio galaxy of comparable power to 3C sources at z$\\sim$1. In consequence, it is one of the most well-studied radio galaxies. Previous work has concentrated on using Cygnus A to test the unified schemes for powerful radio galaxies (\\citealt{ueno94}; \\citealt{ogle97}; \\citealt{young02}) and to search for signs of the impact of the activity on the host galaxy in the form of AGN-driven outflows (\\citealt*{taylor03}). Dynamical studies have also led to estimates of the supermassive black hole mass (\\citealt{tadhunter03}). However, despite the large amount of progress that has been made over the last decade in the understanding of Cygnus A and its host galaxy, relatively little is known about the events that triggered the activity. In this paper we present near-IR spectroscopic observations, of higher resolution than previously published near-IR data (\\citealt{ward91}; \\citealt*{tsr99}; \\citealt{wilman00}), that provide evidence for non-equilibrium gas motions that may have a direct bearing on the triggering of the activity in this poweful AGN. In order to put our observations in context, we start by summarising our current state of knowledge of the geometry and kinematics of the near-nuclear regions of Cygnus A. \\subsection{Geometry and inclination of the near-nuclear regions} HST imaging has revealed an edge-brightened biconical structure on the scale of a few hundred parsecs (\\citealt{jackson98}; \\citealt{tadhunter99}). This structure is most likely a direct result of the nuclear-driven winds hollowing out `ionisation cones' in the kpc-scale dust lane in which the AGN is embedded. Moreover, the images also show evidence for a starburst ring associated with the dust lane (\\citealt{fosbury99}). From these data, the best estimates for the geometry of the biconical structure are opening half-angle $\\theta_{\\frac{1}{2}}$ $\\simeq$ 60\\degr and inclination of the cone axis relative to the line-of-sight i $\\simeq$ 30\\degr with the NW cone oriented towards us and the SE cone oriented away from us (see \\citealt{tadhunter03} for a full discussion). From this geometry we know that we are likely to be viewing Cygnus A very close to the opening angle of the NW cone, but we also know that i $<$ (90 - $\\theta_{\\frac{1}{2}}$) as we do not have a direct view of the quasar nucleus at optical wavelengths. \\subsection{Kinematic components in the near-nuclear regions} The central regions of Cygnus A show considerable complexity in their emission line kinematics. The various kinematic components can be tied to different processes associated with the AGN as follows: \\subsection*{Gravitational motions} {\\bf Kpc-scale disc:} near-IR Keck II spectroscopic data shows evidence for a rotating disc associated with the 3 kpc scale dust lane. The gas is in Keplerian rotation about a stellar core and an unresolved point mass of (2.5 $\\pm$ 0.7) $\\times$ 10$^9$ solar masses. The rotation axis of the gas is aligned at $\\sim$9\\degr to that of the large-scale radio axis (\\citealt{tadhunter03}). \\vspace{2mm} \\noindent{\\bf 300pc-radius disc:} Higher spatial resolution HST/STIS optical data of \\citet{tadhunter03} show evidence for a smaller disc structure. The lines from this disc are significantly broadened throughout (FWHM 500-900 km s$^{-1}$). Although the slit positions closest to the nucleus show evidence for rotation about a central black hole with mass similar to that deduced from the KeckII/NIRSPEC data, there is evidence for deviations from pure circular motions in the NW cone. \\subsection*{Evidence for outflows} {\\bf Emission line outflows:} Intermediate spectral resolution optical data of the [O III]$\\lambda$5007 line (\\citealt{taylor03}) provide evidence for outflowing gas in the ionisation cones. A $\\sim$300 km s$^{-1}$ outflow is detected in the NW cone (blueshifted), whereas a $\\sim$400 km s$^{-1}$ outflow is detected in the SE cone (redshifted). These results are consisitent with gas being driven out of the cones by AGN-induced winds. \\vspace{2mm} \\noindent{\\bf Extreme outflow:} an extreme outflow component is detected in [O III]$\\lambda\\lambda$(5007,4949) emission in the NW cone (\\citealt{tadhunter91}). This component lies on the radio axis itself and is blueshifted by 1300 - 1800 km s$^{-1}$; it is likely to represent gas entrained in the outflowing radio plasma. \\vspace{2mm} \\noindent{\\bf Scattering outflow:} The redshifted [O III]$\\lambda$5007 feature detected in polarised light, in both cones and the nucleus, by \\citet{vanbemmel03} is consistent with the presence of a scattering outflow with velocity in the range 150-450 km s$^{-1}$. \\subsection*{Evidence for inflows} {\\bf Infalling H I:} Two components of H I 21cm absorption are detected in the nuclear regions of Cygnus A (\\citealt{conway95}). In the rest frame, defined by the narrow component of Pa$\\alpha$ in the nuclear aperture (see below), these components have {\\it redshifts} of 227 $\\pm$ 9 km s$^{-1}$ and 48 $\\pm$ 9 km s$^{-1}$. Since these absorption components are detected against the radio core, and are therefore in the foreground, they must be associated with infalling material. \\vspace{2.5mm} \\noindent Although there is now substantial evidence for the gravitational motions and AGN-induced outflows in Cygnus A, the evidence for inflowing material is relatively sparse and rests solely in the H I observations of \\citet{conway95}. However, it is the inflow component that is most likely to provide the strongest clues to the triggering of the activity. Therefore it is important to obtain further information about the kinematics and distribution of the inflowing gas. In this context we note that \\citet{canalizo03} have recently reported the discovery of a red, secondary point source in the nuclear regions, $\\sim$400pc southwest of the radio nucleus. They argue that this secondary nucleus represents the debris of a minor merger that may be fuelling the AGN activity. We assume the cosmological parameters of H$_0$=75 km s$^{-1}$ Mpc$^{-1}$ and q$_0$=0 throughout this paper. For these parameters 1.00 arcsec corresponds to 1.00 kpc at the redshift of Cygnus A. ", "conclusions": "Our near-IR data have further strengthened the case for the existence of complex, non-gravitational kinematics in the core of the nearby radio galaxy Cygnus A. We show several distinct Pa$\\alpha$ components, both on- and off-nucleus, that are consistent with the [O III] components seen in the optical data of \\citet{taylor03}. The molecular hydrogen emission, however, shows markedly different kinematics that are inconsistent with both the Pa$\\alpha$ and optical data. We have found good evidence for the existence of an infalling molecular cloud to the NW of the nucleus and have argued that this lies outside the main ionisation cone and is being excited by the more penetrating hard X-ray emission from the AGN. The redshift of this infalling component is in excellent agreement with that of the H I absorption detected by \\citet{conway95}, suggesting the two phenomena are linked. The fact that the H I absorption is seen on-nucleus, whereas the H$_2$ emission is seen at a projected radius of $\\sim$1.35 kpc, suggests that the infalling material is likely to have a complex and extended spatial structure. We suggest the possibility that the non-circular motions and infalling material in the nuclear regions are linked to a possible merger event (\\citealt{canalizo03}) that may have triggered the AGN activity itself. However, the problem remains that these long-slit data give kinematic information only along the radio axis, whereas a minor merger event could be expected to lead to non-gravitational motion and flows throughout the cones. Future near-IR integral field data that map the structure and kinematics of the H$_2$ emission throughout the cone will reveal if the redhifted emission is associated with the radio jet itself or is more evenly distributed in the near-nuclear regions. Near-IR adaptive optics spectroscopy of the nuclear regions will also be able to clarify the nature of the asymmetrically shifted H$_2$ components seen in our nuclear aperture." }, "0405/astro-ph0405168_arXiv.txt": { "abstract": "The masses of central supermassive black holes in a soft X-ray selected sample of the narrow-line Seyfert 1 galaxies (NLS1s) are estimated by some different methods to test their theoretical models. Apart from the methods using the H$\\beta$ linewidth and the [O III] linewidth, soft X-ray excess as a prominent character of NLS1s is used to estimate the black hole masses. The virial mass derived from the H$\\beta$ linewidth assuming random orbits of broad-line reigns (BLRs) is consistent with that from the soft X-ray bump luminosity for NLS1s but with a larger scatter. The virial black hole masses showed that most of NLS1s are in the super-Eddington accretion state while most of broad-line Seyfert 1 galaxies (BLS1s) are not. We found that the black hole mass estimated from [O III] linewidth is not in agreement with above two methods. Using the Eddington limit relation for the super-Eddington accretion suggested by Wang (2004), we found that there are 16 NLS1s satisfied with this Eddington limit relation. The masses of these 16 NLS1s derived from X-ray luminosity are systematically larger than that from H$\\beta$ linewidth assuming random BLRs orbits. If the mass derived from X-ray luminosity is true, the mean disk inclination to the line of sight in these 16 NLS1s is about $17^{\\circ}$, which provided new support for the pole-on orientation effect in NLS1s. ", "introduction": "Narrow-line Seyfert 1 galaxies (NLS1s) are a peculiar class of active galactic nuclei (AGNs). They are characterized (Osterbrock \\& Pogge 1985): smaller H$\\beta$ FWHM (less than $2000km~s^{-1}$), strong optical Fe II multiplets, the line luminosity ratio of [O~III] 5007$\\AA$ to H$\\beta$ is less than 3, steep soft X-ray excess (Boller et al. 1996), and rapid soft/hard X-ray variability (Leighly 1999; Cheng et al. 2002). A popular model of NLS1 is that they contain less massive black holes, but have higher accretion rates radiating at close Eddington luminosity, namely high Eddington ratios (Pounds et al. 1995; Laor et al. 1997; Mineshige et al. 2000). It has been suggested that NLS1s might be in the early stage of AGNs evolution (Grupe 1996, Grupe et al. 1999, Mathur 2000; Bian \\& Zhao 2003a) and black hole grows fast via a higher fraction of Eddington accretion rate (Sulentic et al. 2000; Boroson 2002). To test this hypothesis we need to estimate the black hole mass of NLS1s. There are several methods to calculate the black hole masses in AGNs: 1) virial mass derived from the H$\\beta$ FWHM and the sizes of broad line regions (BLRs) from the reverberation mapping technique or the empirical size-luminosity formula (Ho 1998; Wandel et al. 1999; Kaspi et al. 2000; Peterson et al. (2000); Vestergaard 2002; Bian \\& Zhao 2003b; 2004); 2) soft X-ray variability (Czerny et al. 2001; Bian \\& Zhao 2003c); 3) the relation between the mass and the bulge stellar velocity dispersion (M-$\\sigma$ relation) or the bulge luminosity (M-$L_{bulge}$ relation)(McLure \\& Dunlop 2001; Tremaine 2002). We are not sure whether above methods can apply to NLS1s for their special properties. Applying reverberation mapping technique to NLS1s is difficult since they are usually less variable in optical band (Shemmer \\& Netzer 2000). Whether the assumption of the random BLRs orbits in NLS1s is suitable or not is a question open to debate (Bian \\& Zhao 2002). The empirical size-luminosity relation is only based on many broad-line Seyfert 1 galaxies (BLS1s) and a few NLS1s (Kaspi et al. 2000). Peterson et al. (2000) suggested that NLS1s and BLS1s follow the same size-luminosity relation. However, BLRs physics in NLS1s is possibly special and different compared with BLS1s (e.g. Boller et al. 1996). Therefore, whether NLS1s follow this relation or not should be confirmed by future observation in NLS1s. The $M-\\sigma$ relation may not apply to NLS1s since they are most likely in the early stage of AGNs evolution, which is intimately related with the black hole growth process (Mathur et al. 2001; Wandel 2002; Lu \\& Yu 2003; Shields et al. 2003; Bian \\& Zhao 2004; Grupe \\& Mathur 2004). As an extreme feature in NLS1s, Soft X-ray excess may most likely be caused by high accretion rate in units of Eddington accretion rate. This character could be used to probe the black hole mass since the photon trapping effect gives a saturated luminosity, namely the luminosity is almost independent to the accretion rate (Wang \\& Zhou 1999; Wang et al. 1999; Ohsuga et al. 2002). Recently Wang \\& Netzer (2003) presented a extreme slim disk with a hot corona to explain the soft X-ray bump in NLS1s and suggested that soft X-ray humps in NLS1s are natural consequences of super-Eddington accretion. They found that the hump X-ray luminosity is weakly dependent on the accretion rate and almost completed determined by black hole mass in their model: \\begin{equation} M_{\\rm BH}=2.8\\times 10^{6}(\\frac{L_{\\rm SX}}{10^{44}\\rm erg~ s^{-1}})\\Msolar \\end{equation} Where $L_{\\rm SX}$ is the soft X-ray luminosity in the flat part of the bump. It provides new method to estimate the black hole masses in NLS1s with super-Eddington accretion rates. In this paper, we compared the results from different methods to estimate the black hole masses in a sample of NLS1s and BLS1s (Grupe et al. 2004). We tried to find which method is suitable to estimate the black hole masses in NLS1s and whether the disk inclinations to the line of sight in NLS1s compared with BLS1s are small or not. All of the cosmological calculations in this paper assume $H_{0}=75 \\rm {~km ~s^ {-1}~Mpc^{-1}}$, $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda} = 0.7$. ", "conclusions": "Different methods are used to estimate the black hole masses in a sample of NLS1s and BLS1s (Grupe et al. 2004). The main conclusions can be summarized as follows: \\begin{itemize} \\item{The mass from the soft X-ray bump luminosity is consistent with that from the H$\\beta$ linewidth for NLS1s. Most of NLS1s are in the super-Eddington accretion state considering $L_{bol}/L_{Edd}>1$. The black hole masses of NLS1s from the H$\\beta$ linewidth and from soft X-ray luminosity are reliable while that from [O III] linewidth are not reliable.} \\item{The mean disk inclination to the line of sight in 16 NLS1s satisfied with equation (5) is about $17^{\\circ}$, which provided new support for the pole-on orientation effect in NLS1s.} \\end{itemize}" }, "0405/astro-ph0405442_arXiv.txt": { "abstract": "We have analyzed galaxy and group-sized dark matter halos formed in a high resolution \\lcdm\\ numerical N-body simulation in order to study the rotation of the triaxial figure, a property in principle independent of the angular momentum of the particles themselves. Such figure rotation may have observational consequences, such as triggering spiral structure in extended gas disks. The orientation of the major axis is compared at 5 late snapshots of the simulation. Halos with significant substructure or that appear otherwise disturbed are excluded from the sample. We detect smooth figure rotation in 278 of the 317 halos in the sample. The pattern speeds follow a log normal distribution centred at $\\Omega_p = 0.148~\\hkmskpc$ with a width of 0.83. These speeds are an order of magnitude smaller than required to explain the spiral structure of galaxies such as \\ntnif. The axis about which the figure rotates aligns very well with the halo minor axis, and also reasonably well with its angular momentum vector. The pattern speed is correlated with the halo spin parameter $\\lambda$, but shows no correlation with the halo mass. The halos with the highest pattern speeds show particularly strong alignment between their angular momentum vectors and their figure rotation axes. The figure rotation is coherent outside 0.12~\\rvir. The measured pattern speed and degree of internal alignment of the figure rotation axis drops in the innermost region of the halo, which may be an artifact of the numerical force softening. The axis ratios show a weak tendency to become more spherical with time. ", "introduction": "The basic method is to identify individual halos in the final $z=0$ snapshot of the simulation, to find their respective progenitors in slightly earlier snapshots, and to measure the rotation of the major axes through their common plane as a function of time. Precisely determining the direction of the axes is crucial and difficult. When merely calculating axial ratios or internal alignment, errors on the order of a few degrees are tolerable. However, if a pattern speed of 1~\\kmskpc, as observed in the halo of \\citetalias{bureau-etal99}, is typical, then a typical halo will only rotate by 4\\degr\\ in between the penultimate and final snapshots of the simulation. Therefore, the axes of a halo must be determined more precisely than this in order for the figure rotation to be detectable. In fact, we should strive for even smaller errors to see if slower-rotating halos exist. It would have been difficult for \\citetalias{pfitzner99} to detect halos rotating much slower than the halo presented in \\citetalias{bureau-etal99}; although the error varies from halo to halo (for reasons discussed in section~\\ref{axis error section}), Figure~5.23 of \\citetalias{pfitzner99} shows that most of his halos had orientation errors of between 8\\degr\\ and 15\\degr, corresponding to a minimum resolvable figure rotation of $\\sim 0.6~\\kmskpc$ for a $2\\sigma$ detection in snapshots spaced 500~Myr apart. A major difficulty in determining the principal axes so precisely is substructure. The orientation of a mass distribution is usually found by calculating the moment of inertia tensor $I_{ij} = \\sum_k m_k r_{k,i} r_{k,j}$, and then diagonalizing $I_{ij}$ to find the principal axes. However, this procedure weights particles by $r^2$. Therefore, substructure near the edge of the halo (or of the subregion of the halo used to calculate the shape) can exert a large influence on the shape of nearly spherical halos, especially if a particular subhalo is part of the calculation in one snapshot but not in another, such as when it has just fallen in. This is particularly problematic because subhalos are preferentially found at large radii \\citep{ghigna-etal00,delucia-etal04,gkg04,gao-etal04}. Moving substructures can also induce a false measurement of figure rotation due to their motion within the main halo at approximately the circular velocity. To mitigate this, we firstly use particles in a spherical region of radius 0.6~\\rvir\\ surrounding the center of the halo, rather than picking the particles from a density dependent ellipsoid as in \\citet{warren-etal92} or \\citet{js02}. We find that those methods allow substructure at one particular radius to influence the overall shape of the ellipsoid from which particles are chosen for the remainder of the calculation, and therefore bias the results even when other measures are adopted to minimize their effect. The choice of a spherical region biases the derived axis ratios toward spherical values, but does not affect the orientation. Secondly, the particles are weighted by $1/r^2$ so that each mass unit contributes equally regardless of radius \\citep{gerhard83}. Both \\citetalias{dubinski92} and \\citetalias{pfitzner99} take similar approaches, but using radii based on ellipsoidal shells. Therefore, we base our analysis on the principal axes of the reduced inertia tensor \\begin{equation}\\label{eq inertia tensor} \\tilde{I}_{ij} = \\sum_k \\frac{ r_{k,i} r_{k,j} }{r_k^2}. \\end{equation} In the majority of halos, the substructure is a small fraction of the total mass, usually less than 5\\%\\ of the total mass within 60\\%\\ of the virial radius \\citep[Figure~8]{delucia-etal04}, so its effect is much reduced. There are still some halos which have not yet relaxed from a recent major merger, in which case the ``substructure'' constitutes a significant fraction of the halo. To find these cases, the reduced inertia tensor is separately calculated enclosing spheres of radius 0.6, 0.4, 0.25, 0.12, and 0.06 times the virial radius to look for deviations as a function of radius (see section~\\ref{5sigma deviations} for details). These radii are always with respect to the $z=0$ value of \\rvir. We find that only halos with at least $4 \\times 10^3$ particles, or masses of at least $\\sim 3 \\times 10^{11}~\\hmsun$ have sufficient resolution for the orientation of the principal axis to be determined at sufficient precision (see Section~\\ref{axis error section}). There are 1432 halos in the $z=0$ snapshot satisfying this criterion, with masses extending up to $2.8 \\times 10^{14}~\\hmsun$. \\subsection{Halo matching}\\label{halo matching section} \\begin{figure} \\plotone{f1.eps} \\caption{\\label{progenitor mass fraction plot} Histogram of the fraction of the final mass that comes from the b096 ($z \\approx 0.05$, solid line) and b090 ($z \\approx 0.12$, dashed line) halo which contributes the most mass.% } \\end{figure} To match up the halos at $z=0$ with their earlier counterparts, we use the individual particle numbers provided by GADGET which are invariant from snapshot to snapshot, and find which halo each particle belongs to in each snapshot. The progenitor of each $z=0$ halo in a given $z>0$ snapshot is the halo that contributes $\\ge 90\\%$ of the final halo mass. Sometimes no such halo exists; in these cases, the halo has only just formed or underwent a major merger and so is not useful for our purposes. Figure~\\ref{progenitor mass fraction plot} shows a histogram of the fraction of the final halo mass that comes from the b096 ($z \\approx 0.05$) halo which contributes the most mass. There are also some cases where two nearby objects are identified as a single halo in an earlier snapshot but as distinct objects in the final snapshot. We therefore impose the additional constraint that the mass contributed to the final halo must also be $\\ge 90\\%$ of the progenitor's mass. In the longer time between the earliest snapshot b090 and the final snapshot b102, a halo typically accretes a greater fraction of its mass, and so a more liberal cut of 85\\%\\ is used for this snapshot (see the dashed histogram in Figure~\\ref{progenitor mass fraction plot}). 492 of the halos that satisfied the mass cut did not have a progenitor which satisfied these criteria in at least one of the $z>0$ snapshots and so were eliminated from the analysis, leaving a sample of 940 matched halos. \\subsection{Error in axis orientation}\\label{axis error section} There are two sources of errors that enter into the determination of the axes: how well the principal axes of the particle distribution can be determined, and whether that particle distribution has a smooth triaxial figure. Here we estimate the error assuming that it is not biased by substructure. The halos for which this assumption does not hold will become apparent later in the calculation. For a smooth triaxial ellipsoid represented by $N$ particles, the error is a function of $N$ and of the intrinsic shape: as the axis ratio $b/a$ or $c/b$ approaches unity, the axes become degenerate. To quantify this, we have performed a bootstrap analysis of the particles in a sphere of radius 0.6~\\rvir\\ of each $z=0$ halo \\citep{heyl-etal94}. If the sphere contains $N$ particles then we resample the structure by randomly selecting $N$ particles from that set allowing for duplication and determine the axes from this bootstrap set. We do this 100 times for each halo. The dispersion of these estimates around the calculated axis is taken formally as the ``$1\\sigma$'' angular error, and is labelled \\thboot. \\begin{figure} \\plotone{f2-colour.eps} \\caption{\\label{major axis error vs N} Angular bootstrap error \\thboot\\ as a function of the number of particles $N$ within the central 0.6~\\rvir\\ of each halo. Points are the cosmological halos, and asterisks are randomly sampled smooth NFW halos. \\textit{(Top):} Angular error \\thboot. The solid line is the fit \\therrN\\ from equation~(\\ref{theta err N}). \\textit{(Middle):} Ratio between the angular error and the error expected for the halo given its axis ratio $b/a$, i.e. $\\thboot / \\therrba$. The solid line is \\therrN\\ from equation~(\\ref{theta err N}) renormalized to the typical error of 0.02~radians. \\textit{(Bottom):} Ratio between the angular error and the analytic estimate \\therr\\ from equation~(\\ref{predicted error}). } \\end{figure} \\begin{figure} \\plotone{f3-colour.eps} \\caption{\\label{major axis error vs b/a} Angular bootstrap error \\thboot\\ as a function of the axis ratio $b/a$ of each halo. Points are the cosmological halos, and asterisks are randomly sampled smooth NFW halos. \\textit{(Top):} Angular error \\thboot. The solid line is the fit \\therrba\\ from equation~(\\ref{theta err b/a}). \\textit{(Middle):} Ratio between the angular error and the error expected for the halo given the number of particles $N$, i.e. $\\thboot / \\therrN$. The solid line is \\therrba\\ from equation~(\\ref{theta err b/a}) renormalized to the typical error of 0.02~radians. \\textit{(Bottom):} Ratio between the angular error and the analytic estimate \\therr\\ from equation~(\\ref{predicted error}). } \\end{figure} As expected, the two important parameters are $N$ and the axis ratio. We focus here on the major axis, for which the important axis ratio is $b/a$. The top panels of Figures~\\ref{major axis error vs N} and~\\ref{major axis error vs b/a} show the dependence of the bootstrap error on $N$ and $b/a$ respectively for all halos with $M > 10^{11}~\\hmsun$. The solid lines are empirical fits: \\begin{equation}\\label{theta err N} \\therrN = \\frac{2}{\\sqrt{N}}, \\end{equation} and \\begin{equation}\\label{theta err b/a} \\therrba = 0.005 \\frac{\\sqrt{b/a}}{1 - b/a}. \\end{equation} The form of equation~(\\ref{theta err N}) is not surprising; if a smooth halo was randomly sampled, we would expect the errors to be Poissonian with an $N^{-1/2}$ dependence. However, the cosmological halos are not randomly sampled. Individual particles ``know'' where the other particles are, because they have acquired their positions by reacting in the potential defined by those other particles. Therefore, the errors may be less than expected from a randomly sampled halo. To test this, we construct a series of smooth prolate NFW halos \\citep*{nfw96} with $b/a$ axis ratios ranging from 0.5~to 0.9, randomly sampled with between $3\\times 10^3$ and $3\\times 10^5$ particles, and perform the bootstrap analysis identically for each of these halos as for the cosmological halos. Because the method for calculating axis ratios outlined in Section~\\ref{methodology intro section} biases axis ratios toward spherical, the recovered $b/a$ of these randomly sampled halos is larger than the input value, and ranges from 0.65~to 0.95. The errors for these randomly sampled smooth halos are shown as asterisks in Figures~\\ref{major axis error vs N} and~\\ref{major axis error vs b/a}. The top panel of Figure~\\ref{major axis error vs N} shows a rise in the dispersion of the error for $N\\lesssim 4000$, with many halos having errors greater than the 0.1~radians necessary to detect the figure rotation of the halo presented in \\citetalias{bureau-etal99}. Therefore, we only use halos with $N>4000$. The bootstrap error appears to be completely determined by $N$ and $b/a$. The residuals of \\thboot\\ with respect to \\therrN\\ are due to \\therrba\\ and vice versa. This is shown in the middle panels of Figures~\\ref{major axis error vs N} and \\ref{major axis error vs b/a}. In the middle panel of Figure~\\ref{major axis error vs N} we have divided out the dependence of \\thboot\\ on the axis ratio, making apparent an extremely tight relation between the residual and $N$, while in the middle panel of Figure~\\ref{major axis error vs b/a} we have divided out the dependence of \\thboot\\ on $N$, showing the equally tight relation between the residual and $b/a$. It is apparent from comparing the points and asterisks that the errors in the cosmological halos are slightly smaller than for randomly sampled smooth halos. Combining equations~(\\ref{theta err N}) and~(\\ref{theta err b/a}), and noting that the typical error is $\\thboot \\approx 0.02$~radians, we find the bootstrap error is well fit by \\begin{equation}\\label{predicted error} \\therr = \\frac{1}{2 \\sqrt{N}} \\frac{\\sqrt{b/a}}{1 - b/a}. \\end{equation} The bottom panels of Figures~\\ref{major axis error vs N} and~\\ref{major axis error vs b/a} show the residual ratio between the bootstrap error \\thboot\\ and the analytic estimate \\therr. The vast majority of points lie between 0.8~and 1.0, indicating that \\therr\\ overestimates the error by $\\sim 10\\%$. Equation~(\\ref{predicted error}) breaks down as $b/a$ approaches unity; these halos are nearly oblate and so do not have well-defined major axes. It also becomes inaccurate at very low $b/a$ due to low-mass poorly-resolved halos. Even in these cases, the error estimate is conservative, but to be safe we have eliminated axes with $b/a < 0.35$ or $b/a > 0.95$ from the subsequent analysis, regardless of the nominal error. The randomly-sampled smooth halos follow equation~(\\ref{predicted error}) extremely well, so the non-Poissonianity of the sampling in simulated halos reduces the errors by 10\\%. Calculating the bootstraps is computationally expensive, so equation~(\\ref{predicted error}) is used for the error in all further computation. Because this estimate is expected to be correct for smooth ellipsoids, cases where the error is anomalous are indications of residual substructure. \\subsection{Figure rotation}\\label{figure rotation methodology section} \\begin{figure} \\plotone{f4.eps} \\caption{\\label{plane diagram}% Diagram that demonstrates how we fit a plane to the major axis measurements at all snapshots (thick lines) and then find the increase of phase $\\phi$ as a function of time. The figure rotation axis is perpendicular to the best fit plane, and defined such that $\\phi$ increases around it counter-clockwise with time.} \\end{figure} Ideally one would fit the figure rotation by comparing the orientation of the axis at each snapshot to that of a unit vector rotating uniformly along a great circle, and minimize the $\\chi^2$ to find the best fit uniform great circle trajectory. However, this requires minimizing a non-linear function in a four-dimensional parameter space, a non-trivial task. We adopt the simpler and numerically more robust method of solving for the plane $z = ax + by$ that fits the major axis measurements of the halo best at all timesteps, assuming the error is negligible. The change of the phase of the axes in this plane as a function of time are then fit by linear regression. A schematic diagram of this process is shown in Figure~\\ref{plane diagram}. The degree to which the axes are consistent with lying in a plane is checked by calculating the out-of-plane $\\chi^2$: \\begin{equation}\\label{out of plane chi2} \\chi^2_{\\mathrm{oop}} \\equiv \\frac{1}{\\nu} \\sum_i \\frac{ \\Delta\\theta_i^2 } {{\\therr}_i^2}, \\end{equation} where $\\nu$ is the number of degrees of freedom and $\\Delta\\theta_i$ is the minimum angular distance between the major axis at timestep $i$ and the great circle defined by the best fit plane. Because the axes have reflection symmetry, it is impossible to measure a change in phase of more than $\\pi/2$. The phases are adjusted by units of $\\pi$ such that the difference in phase between adjacent snapshots is always less than $\\pi/2$. If the figure were truly rotating by 90\\degr\\ or more in between the snapshots, it would be impossible to accurately measure this rotation since the angular frequency would be larger than the Nyquist frequency of our sampling rate. Any faster pattern speeds would be aliased to lower angular velocities, with an aliased angular velocity of $\\Omega_{\\mathrm{Nyq}} - (\\Omega_p - \\Omega_{\\mathrm{Nyq}})$, where $\\Omega_p$ is the intrinsic angular velocity of the pattern and $\\Omega_{\\mathrm{Nyq}}$ is the Nyquist frequency. For snapshots spaced $500~h^{-1}~\\mathrm{Myr}$ apart, the maximum time between the snapshots we analyze, the maximum detectable angular velocity is 3.8~\\hkmskpc. We do not expect the figure to change so dramatically as we have excluded major mergers. However, this can be checked \\textit{post facto} by checking whether the distribution of measured angular velocities extends up to the Nyquist frequency; if so, then there are likely even more rapidly rotating figures whose angular frequency is aliased into the detectable range, fooling us into thinking they are rotating slower. If the measured distribution does not extend to the Nyquist frequency, then it is unlikely that there are any figures rotating too rapidly to be detected (see Section~\\ref{results section}). The best fit linear relation for the phase as a function of time is found by linear regression. Because the component of an isotropic angular error projected onto a plane is half of the isotropic error, we divide the error of equation~(\\ref{predicted error}) by two before we perform the regression. The slope of the linear fit gives the pattern speed $\\Omega_p$ of the figure rotation. The error is the $1\\sigma$ limit on the slope. \\begin{figure} \\plotone{f5.eps} \\caption{\\label{subhalo fraction} Log-weighted projected density of 4 halos with a range of subhalo fractions $f_s$. The subhalo fractions are 0.166 (top-left), 0.065 (top-right), 0.045 (bottom-left), and 0.016 (bottom-right). Axes are in units of $h^{-1}~\\mathrm{kpc}$ from the halo center. All halos have masses in the range $2$ -- $3 \\times 10^{12}~\\hmsun$. } \\end{figure} \\label{5sigma deviations}% Once we have calculated the pattern speed for each halo, we can eliminate the cases where substructure has severely impacted the results. In these cases, the signal is dominated by a large subhalo at a particular radius, so the derived pattern speed will be significantly different when the sphere is large enough to include the subhalo from when the subhalo is outside the sphere. We have calculated the pattern speed using enclosing spheres of radius 0.6, 0.4, 0.25, 0.12, and 0.06 of the virial radius. The fraction of mass in subhalos can be estimated via the change in the pattern speed $\\Omega_p$ at adjacent radii. Because the reduced inertia tensor is mass-weighted, the figure rotation of a sphere with a smooth component rotating at $\\Omega_{p,\\mathrm{smooth}}$ plus a subhalo containing a fraction $f_s$ of the total mass moving at the circular velocity $v_c$ at radius $R~\\rvir$ is approximately \\begin{equation}\\label{omegap for a subhalo} \\Omega_p \\approx (1-f_s) \\Omega_{p,\\mathrm{smooth}} + f_s \\frac{v_c}{R~\\rvir}, \\end{equation} where the difference due to the presence of the subhalo is \\begin{equation}\\label{delta omegap = fs vc / r} \\Delta \\Omega_p = f_s \\frac{v_c}{R~\\rvir} = f_s \\sqrt{ \\frac{ G M(0.95$ discussed in Section~\\ref{axis error section}. We also eliminate cases where the reduced $\\chi^2$ from the linear fit of phase versus time indicates that the intrinsic error of the direction determination is much lower than suggested by equation~(\\ref{predicted error}), indicating that the model of the halo as a smooth ellipsoid is violated (10 halos with $\\chi^2_{\\nu} < 0.1$), and those cases where the phase does not evolve linearly with time (134 halos with $\\chi^2_{\\nu} > 10$). Finally, we eliminate halos where the axes do not lie on a common plane, i.e. the 32 halos where $\\chi^2_{\\mathrm{oop}} > 10$. Therefore, the final sample consists of 317~halos. \\begin{figure} \\plotone{f6-small.eps} \\caption{\\textit{(Upper-left five panels):} Projection onto the best fit plane of the inner 0.6~\\rvir\\ of a sample halo at the five snapshots we analyze. Axes are in units of $h^{-1}~\\mathrm{kpc}$ from the halo center. From left to right, top to bottom, the snapshots are at 1108, 496, 296, 98, and 0~\\hmyr\\ before $z=0$. The solid line is the major axis, which rotates counterclockwise by 20\\degr\\ from beginning to end. \\textit{(Bottom-right):} Phase of the major axis in the rotational plane of the sample halo. The zero point is arbitrary, but identical in all snapshots. The solid line is the linear fit, with a slope of $0.33~h~\\mathrm{radians~Gyr^{-1}}$. \\label{sample halo}} \\end{figure} A sample halo is shown in the first five panels of Figure~\\ref{sample halo}. It was chosen randomly from the halos with relatively low errors and typical pattern speeds. It has a mass of $1.9\\times 10^{12}~\\hmsun$, and a pattern that rotates at $0.32\\pm 0.01~\\hkmskpc$. It has a spin parameter $\\lambda = 0.047$, and axis ratios of $b/a=0.86$ and $c/a=0.77$ at $z=0$. The derived substructure fraction is $f_s=0.045$, and the out-of-plane $\\chi^2_{\\mathrm{oop}} = 8.5$. The solid line shows the measured major axis in each snapshot, which rotates counterclockwise in this projection. The phase of its figure rotation as a function of time is shown in the bottom-right panel of Figure~\\ref{sample halo}. The zero point is arbitrary, but is consistent from snapshot to snapshot. The linear fit is also shown, which has a reduced $\\chi^2$ of 2.9. ", "conclusions": "We have detected rotation of the orientation of the major axis in most undisturbed halos of a \\lcdm\\ cosmological simulation. The axis around which the figure rotates is very well aligned with the minor axis in most cases. It is also usually well aligned with the angular momentum vector. The distribution of pattern speeds is well fit by a log normal distribution, \\begin{equation} P(\\Omega_p) = \\frac{1}{\\Omega_p \\sigma \\sqrt{2\\pi}} \\exp\\left(-\\frac{\\ln^2(\\Omega_p/\\Omega_{p_0})} {2\\sigma^2} \\right), \\end{equation} with $\\Omega_{p_0}=0.148~\\hkmskpc$ and $\\sigma=0.83$. The pattern speed $\\Omega_p$ is correlated with spin parameter $\\lambda$. The median pattern speed rises from 0.12~\\hkmskpc\\ for halos with $\\lambda < 0.02$ to 0.44~\\hkmskpc\\ for halos with $\\lambda > 0.06$, with a spread of almost an order of magnitude around this median at a given value of $\\lambda$. The 11\\%\\ of halos in the sample with the highest pattern speeds, $\\Omega_p > 0.4~\\hkmskpc$, not only have large spin parameters, but also show particularly strong alignment between their figure rotation axes and their angular momentum vectors. There is no obvious correlation of the figure rotation properties with mass. The pattern speed and figure rotation axis is coherent in the outer regions of the halo. Within 0.12~\\rvir, the pattern speed drops, particularly for those halos with fast figure rotation, and the internal alignment of the figure rotation axis deteriorates. This is probably an artifact of the numerical force softening. \\citetalias{bureau-etal99} hypothesized that the spiral structure in \\ntnif\\ is due to figure rotation of a triaxial halo. The required pattern speed of $7\\pm1~\\kmskpc$ \\citepalias{mb03} is much higher than the pattern speeds seen in the simulated halos, and is estimated to have a probability of $5 \\times 10^{-7}$. We therefore conclude that the figure rotation of undisturbed \\lcdm\\ halos is not able to produce this spiral structure. Halos with large values of $\\lambda$ tend to have more substructure \\citep{be87}, so there is a deficiency of halos with very high $\\lambda$ in our sample. Because $\\Omega_p$ correlates with $\\lambda$, we cannot exclude the possibility that there exist halos with very high $\\lambda$ whose figures rotate sufficiently quickly. However, halos with such high $\\lambda$ are themselves very rare \\citepalias{mb03}, and if such halos fall out of our sample due to the presence of strong substructure, the effects of the substructure on the gas disk of \\ntnif\\ would be of more concern than the slow rotation of the halo figure, a possibility \\citetalias{bureau-etal99} rule out due to the lack of any plausible companion in the vicinity. More generally, \\citet{bf02} found very weak if any enhancement of spiral structure in disk simulations with triaxial figures rotating at 0.77~\\kmskpc, a value similar to the highest pattern speed seen in our sample. Therefore, it is unlikely that triaxial figure rotation can be detected by looking for spiral structure in extended gas disks. We have found that the axis ratios of undisturbed halos tend to become more spherical with time, with median fractional increases in the $b/a$ and $c/a$ axis ratios of $\\approx 0.009~h~\\mathrm{Gyr^{-1}}$. The distributions of $(\\dot{b/a})/(b/a)$ and $(\\dot{c/a})/(c/a)$ are relatively wide, with standard deviations of $\\approx 0.03~h~\\mathrm{Gyr^{-1}}$. A few outliers have axis ratios that change quite significantly over the span of 1~Gyr. The rate of change of the axis ratios is not correlated with any other halo property." }, "0405/astro-ph0405218_arXiv.txt": { "abstract": "We report on a search for electro-magnetic and/or hadronic showers (cascades) induced by high energy neutrinos in the data collected with the \\mbox{AMANDA II} detector during the year 2000. The observed event rates are consistent with the expectations for atmospheric neutrinos and muons. We place upper limits on a diffuse flux of extraterrestrial electron, tau and muon neutrinos. A flux of neutrinos with a spectrum $\\Phi \\propto E^{-2}$ which consists of an equal mix of all flavors, is limited to $E^2\\Phi(E)=8.6 \\times 10^{-7} \\stdunit$ at a 90\\% confidence level for a neutrino energy range 50~TeV to 5~PeV. We present bounds for specific extraterrestrial neutrino flux predictions. Several of these models are ruled out. ", "introduction": "The existence of high-energy extraterrestrial neutrinos is suggested by the observation of high-energy cosmic rays and gamma rays. Observation of neutrinos could shed light on the production and acceleration mechanisms of cosmic-rays, which for energies above the ``knee'' ($10^{15} \\, \\mathrm{eV}$) remain not understood. Cosmic rays are thought to be accelerated at the shock fronts of galactic objects like supernova remnants, micro-quasars, and in extragalactic sources such as the cores and jets of active galactic nuclei (AGN) \\cite{halzen02}. High energy protons accelerated in these objects may collide with the gas and radiation surrounding the acceleration region, or with matter or radiation between the source and the Earth. Charged pions, produced in the interaction, decay into highly energetic muon neutrinos and muons which further decay into electron neutrinos. Fermi acceleration of charged particles in magnetic shocks naturally leads to power-law spectra, $E^{-\\alpha}$, where $\\alpha$ is typically close to 2. Hence, the spectrum of astrophysical neutrinos is harder than the spectrum of atmospheric neutrinos ($\\sim E^{-3.7}$) potentially allowing to distinguish the origin of the flux (see for example \\cite{lm}). For a generic astrophysical neutrino source, one expects a ratio of neutrino fluxes $\\Phi_{\\nu_e}:\\Phi_{\\nu_\\mu}:\\Phi_{\\nu_\\tau}\\approx 1:2:0$. Due to neutrino vacuum oscillations this ratio changes to $\\Phi_{\\nu_e}:\\Phi_{\\nu_\\mu}:\\Phi_{\\nu_\\tau} \\approx1:1:1$ by the time the neutrinos reach the Earth. Recently a search with the AMANDA detector was reported \\cite{gary}, resulting in the most restrictive upper limit on the diffuse flux of muon neutrinos (in the energy range 6 to 1000 TeV). Clearly, a high sensitivity to neutrinos of all neutrino flavors is desirable. The present paper reports on a search for a diffuse flux of neutrinos of all flavors performed using neutrino-induced cascades in AMANDA. ", "conclusions": "We have presented experimental limits on diffuse extragalactic neutrino fluxes. We find no evidence for neutrino-induced cascades above the backgrounds expected from atmospheric neutrinos and muons. In the energy range from 50~TeV to 5~PeV, the presented limits on the diffuse flux are currently the most restrictive. We have compared our results to several model predictions for extragalactic neutrino fluxes and several of these models can be excluded. Results from the first phase of AMANDA, the 10-string sub-detector AMANDA-B10, have been reported in \\cite{b10} and an update to the analysis was presented above. Compared to AMANDA-B10, the analysis presented here has a nearly ten times larger sensitivity, mainly achieved through using the larger volume of AMANDA-II and by extending the search to neutrinos from all neutrino directions. The limits presented here are also more than a factor of two below the AMANDA-B10 limit obtained by searching for neutrino-induced muons \\cite{gary} and roughly as sensitive as the extension of that search using AMANDA-II 2000 data \\cite{gary_icrc03}. (Assuming a neutrino flavor ratio of 1:1:1, the numerical limits on the flux of neutrinos of a specific flavor (e.q. $\\nu_\\mu$) reported in the literature are 1/3 of the limits on the total flux of neutrinos.) The limits obtained from a search for cascade-like events by the Baikal collaboration \\cite{baikal} are about 50\\% less restrictive than the limits presented here. With the present analysis one obtains a large sensitivity to astrophysical neutrinos of all flavors and in particular to electron and tau neutrinos. Hence, given the large sensitivity to muon neutrinos of other search channels, AMANDA can be considered an efficient all-flavor neutrino detector." }, "0405/nucl-th0405055_arXiv.txt": { "abstract": "We calculate the correlation functions needed to describe the linear response of superfluid matter, and go on to calculate the differential cross section for neutral-current neutrino scattering in superfluid neutron matter and in color-flavor locked quark matter (CFL). We report the first calculation of scattering rates that includes neutrino interactions with both pair-breaking excitations and low-lying collective excitations (Goldstone modes). Our results apply both above and below the critical temperature, allowing use in simulations of neutrino transport in supernovae and neutron stars. ", "introduction": "The ground state of QCD at large densities is color superconducting quark matter \\cite{Barrois:1977xd, Bailin:1984bm, Alford:1998zt,Rapp:1998zu}. When the effects of quark masses can be neglected, three flavor quark matter will be in a particularly symmetric phase called the Color-Flavor Locked (CFL) phase, with BCS pairing involving all nine quarks \\cite{Alford:1998mk}. In this phase the fermion excitation spectrum has a gap $\\Delta$, and model calculations indicate that $\\Delta \\approx 10 - 100$ MeV for quark chemical potential $\\mu \\approx 300 - 500$ MeV. This phase breaks the ${SU(3)_{\\rm color}} \\times SU(3)_L \\times SU(3)_R \\times U(1)_B$ symmetry of QCD down to the global diagonal $SU(3)$ symmetry. The lightest excitations are an octet of pseudo-Goldstone bosons and a true Goldstone boson associated with the breaking of the global $U(1)_B$ symmetry. At some densities, however, the strange quark mass may induce an appreciable stress on the symmetric CFL state, and less symmetric phases may be possible. One possibility is the CFLK$^0$ phase, which exhibits a Bose condensate of $K^0$ in addition to the diquark condensate of the CFL phase--this phase breaks hypercharge and isospin symmetries \\cite{Bedaque:2001je, Kaplan:2001qk}. Another possibility is the LOFF phase, which exhibits crystalline color superconductivity--the diquark condensate varies periodically in space, breaking translation and rotation symmetries \\cite{Alford:2000ze, Bowers:2001ip, Leibovich:2001xr, Kundu:2001tt, Casalbuoni:2001ha, Casalbuoni:2002hr, Casalbuoni:2002my,Casalbuoni:2002pa}. Most recently, a superconducting phase of three-flavor quark matter with non-trivial gapless fermionic excitations has been suggested \\cite{Alford:2003fq}. The densities at which color superconducting quark matter exists could be attained in compact ``neutron'' stars or core-collapse supernovae. It is therefore important to explore the impact of color superconductivity on observable aspects of these astrophysical phenomena. Investigations to that end have included studies of magnetic properties of neutron stars \\cite{Alford:1999pb} and the equation of state of dense matter \\cite{Alford:2002rj, Baldo:2002ju, Banik:2002kc,Shovkovy:2003ce,Buballa:2003et,Ruster:2003zh}. The interactions of neutrinos with superconducting quark matter has also been explored \\cite{Carter:2000xf, Jaikumar:2001hq, Reddy:2003ap}. The emission of neutrinos from CFL during the long-term cooling epoch was studied in Ref. \\cite{Jaikumar:2002vg}. During this epoch the temperature is $T \\lesssim 10^{10}$ K, and most of the excitations of CFL have small number densities. As a result, neutrino emission is highly suppressed. Complementing Ref.~\\cite{Jaikumar:2002vg} is Ref.~\\cite{Reddy:2002xc}, which studied the emission of neutrinos from CFL in a young, proto-neutron star. During this epoch, the temperature is $T \\approx 10^{11}$ K, and the number densities of the excitations are significant. The authors of Ref.~\\cite{Reddy:2002xc} only studied the interactions of neutrinos with the Goldstone bosons, using the effective field theory relevant at energies small compared to the gap. But one might expect the fermionic excitations to become relevant at $T \\approx 10^{11}$ K, since these temperatures are close to the critical temperature of CFL, and the gap will be suppressed. These are the circumstances we analyze in this paper. We study neutral-current neutrino scattering in quark matter, above and below the critical temperature for CFL. We do this by calculating the quark polarization tensor, including the effects of pair-breaking, fermionic excitations. We also include the effects of the collective, bosonic excitation associated with the breaking of $U(1)_B$ by using the Random Phase Approximation (RPA) to build this mode out of microscopic quark-quark interactions. This excitation has the largest contribution of any of the bosonic modes. We begin setting up the calculation in Section II. But before proceeding to study CFL, we consider the case of non-relativistic fermionic modes, to make contact with previous work and to motivate the use of RPA, arguing that consistency with current conversation requires the inclusion of the collective mode--we do this in Section III, where we also go on the calculate the differential cross section for neutrino scattering in superfluid neutron matter. In Section IV we compute the medium polarization tensor for a one component relativistic superfluid. In Section V we compute the medium polarization tensor for CFL, and go on to calculate the differential cross section. We conclude in Section VI with reflections on this work. Further details of the calculation for the one-component, relativistic superfluid can be found in Appendix A. Further details of the calculation for CFL can be found in Appendix B. ", "conclusions": "We have calculated the differential cross section for neutral-current neutrino scattering in superfluid neutron matter, plotted in Fig.~\\ref{nrdiffcs} and in color-flavor locked quark matter, plotted in Fig.~\\ref{cfl_dsigma_Fig}, under conditions relevant to proto-neutron stars. Our results apply above and below the critical temperature. In both of these regimes our model for the interaction in the medium includes the dominant contribution to the cross section. Above the critical temperature, this comes from the fermionic excitations, which become the pair-breaking excitations below the critical temperature. Well below the critical temperature the dominant contribution comes from the massless bosonic mode associated with the breaking of $U(1)$. Although we presented results results for scattering cross sections with space-like kinematics, our polarization functions extend into the time-like region where pair-breaking (and recombination) is the dominant source of the response. These could be employed in calculations of the neutrino emissivity. In particular, we have demonstrated the importance of vertex corrections in these region. Our results suggest that earlier calculations of the neutrino emissivity from the pair recominbation process in superfluid neutron matter \\cite{Flowers:1976} and quark matter \\cite{Jaikumar:2001hq}, which ignore these vertex (RPA) corrections, need to be revised. The analysis presented here is based on mean-field and the random phase approximation (RPA). Its validity is restricted to weak coupling, $\\Delta \\ll \\mu$. For strong coupling we can expect the response to differ quantitatively. In particular, screening corrections which were discussed earlier could be significant. Another drawback is our use of simplified interactions to describe superfluid neutron and quark matter. Our focus was to explore the role of superfluidity and this motivated our choice of a simple zero-range s-wave interaction. In reality, the nucleon-nucleon and quark-quark interactions are more complex. These will induce additional correlations which will affect both the gap equation and the response. (For a recent review on the role of strong interaction correlations on neutrino opacities see Ref.~\\cite{Burrows:2004vq}.) Given the non-perturbative nature of these corrections it is difficult to foresee how large they may be. Although our results are valid both above and below the critical temperature, we have implicitly assumed that the transition is a BCS-like second order transition. Several caveats must be borne in mind when using our results near $T_c$. In the real system this transition may be first order either due to gauge field fluctuations \\cite{Giannakis:2001wz} or stresses such as the strange quark mass and electric charge neutrality \\cite{Steiner:2002gx}. Also, fluctuations of the magnitude of the order parameter dominate in a region called the Ginsburg region around $T_c$. In strong coupling, the size of this region could be significant fraction of $T_c$ \\cite{Voskresensky:2003wd}. Our approach captures some of these fluctuations through RPA. Nonetheless, a Landau-Ginsburg approach - an effective theory for $|\\Delta|$, is more appropriate in this regime \\cite{Iida:2000ha}. In particular, there are precursor fluctuations just above $T_c$, which are not included in our response, that may be relevant \\cite{Kitazawa:2001ft}. These effects are currently under investigation and will be reported elsewhere. Our primary goal was to provide expressions for the differential cross sections that could be used in simulations of the early thermal evolution of neutron stars born in the aftermath of a supernova explosion. The microphysics of neutrino scattering affects the rate of diffusion, which in turn affects macroscopic observables such as the cooling rate and the neutrino emission from core-collapse supernovae. Our results, which extends both to the low and high temperature regions, are well suited for use in simulations of core-collapse supernovae and early thermal evolution of neutron stars. \\vskip0.5in \\centerline{\\bf Acknowledgments} We would like to thank B. Fore, M. Forbes, K. Fukushima, C. Kouvaris, K. Rajagopal and G. Rupak for helpful discussions. J.K. would also like to thank the Institute for Fundamental Theory at the University of Florida for its hospitality. The research of J.K. is supported by the Department of Energy under cooperative research agreement DE-FC02-94ER40818. The research of S. R is supported by the Department of Energy under contract W-7405-ENG-36." }, "0405/astro-ph0405611_arXiv.txt": { "abstract": "We model how repeated supernova explosions in high-redshift dwarf starburst galaxies drive superbubbles and winds out of the galaxies. We compute the efficiencies of metal and mass ejection and energy transport from the galactic potentials, including the effect of cosmological infall of external gas. The starburst bubbles quickly blow out of small, high-redshift, galactic disks, but must compete with the ram pressure of the infalling gas to escape into intergalactic space. We show that the assumed efficiency of the star formation rate dominates the bubble evolution and the metal, mass, and energy feedback efficiencies. With star formation efficiency $f_{*}=0.01$, the ram pressure of infall can confine the bubbles around high-redshift dwarf galaxies with circular velocities $v_c\\ga52$ km s$^{-1}$. We can expect high metal and mass ejection efficiencies, and moderate energy transport efficiencies in halos with $v_c\\approx30-50$ km s$^{-1}$ and $f_*\\approx0.01$ as well as in halos with $v_c\\approx100$ km s$^{-1}$ and $f_*\\gg 0.01$. Such haloes collapse successively from 1--2$\\sigma$ peaks in $\\Lambda$CDM Gaussian density perturbations as time progresses. These dwarf galaxies can probably enrich low and high-density regions of intergalactic space with metals to $10^{-3}$--$10^{-2}$ Z$_{\\odot}$ as they collapse at $z\\approx8$ and $z\\la5$ respectively. They also may be able to provide adequate turbulent energy to prevent the collapse of other nearby halos, as well as to significantly broaden Lyman-$\\alpha$ absorption lines to $v_{rms}\\approx20$--40 km s$^{-1}$. We compute the timescales for the next starbursts if gas freely falls back after a starburst, and find that, for star formation efficiencies as low as $f_{*} \\la 0.01$, the next starburst should occur in less than half the Hubble time at the collapse redshift. This suggests that episodic star formation may be ubiquitous in dwarf galaxies. ", "introduction": "\\label{sec:intro} \\subsection{Overcooling and Angular Momentum Problems} The theory of hierarchical structure formation in a universe dominated by cold dark matter (CDM) predicts that galaxies assembled from mergers of dark matter halos that gravitationally collapsed from primordial fluctuations. Within the same framework, the first luminous structures to form have sub-galactic masses, given by the Jeans mass of the gas. This in turn is determined by the cooling mechanisms dominant during each cosmological era, since CDM models do not define a lower limit to the scale of inhomogeneities in the overall mass distribution. Protogalactic gas clouds of mass $M_h\\sim10^{9}$~M$_{\\odot}$ begin to collapse as $\\sim3\\sigma$ peaks around redshift $z\\sim10$ and become dwarf galaxies. The hierarchical scenario predicts the formation of numerous dwarf-sized halos at high redshift. In these halos, the dissipative collapse of gas is very efficient, since the cooling time $\\tau\\propto\\rho^{-1}\\propto(1+z)^{-3}$, where $\\rho$ is the mass density. This is often referred to as the ``overcooling problem'' (White \\& Rees 1978; White \\& Frenk 1991). Overcooling leads to a prediction of far more dwarf galaxies than observed, as well as to an angular momentum problem. When dense, cooled, protogalactic gas clouds or dwarf galaxies interact within a larger system such as the halo of a present-day disk galaxy, dynamical friction transfers their orbital angular momentum to the surrounding dark matter halo. The result is a galaxy with too much mass, rotating too slowly, with angular momentum deficient by a factor of $\\sim25$ relative to that observed (e.g.\\ Katz \\& Gunn 1991; Navarro \\& Benz 1991; Navarro \\& White 1994). Both of these problems can be resolved if a mechanism can be found to keep the gas diffuse until the peak of present-day disk ($L_*$) galaxy formation. Supernova feedback from young stars is one commonly invoked mechanism to do this. A number of cosmological hydrodynamic simulations have demonstrated that stellar feedback plays a major role in the formation of galaxies (Katz 1992; Navarro \\& White 1993; Yepes et al.\\ 1997; Gerritsen \\& Icke 1997; Hultman \\& Pharasyn 1999; Scannapieco et al.\\ 2001; Thacker \\& Couchman 2001; Sommer-Larsen et al.\\ 2002; Springel \\& Hernquist 2003). However, the qualitative outcomes of simulations vary drastically with the feedback scheme applied. Recently, smoothed particle hydrodynamic (SPH) simulations by Scannapieco et al.\\ (2001) and Springel \\& Hernquist (2003) show that overcooling in low-mass halos is significantly suppressed by the galactic outflows that they model, since the outflows slow down the accretion of gas in the halos and also strip baryons from neighboring halos (\"baryonic stripping\"). The latest high-resolution ($\\sim10^{7}h^{-1}$M$_{\\odot}$ per gas particle) SPH simulations by Thacker and Couchman (2001) and Sommer-Larsen et al.\\ (2002) show that their treatment of thermal feedback with suppressed cooling succeeds in reproducing a present-day disk galaxy with angular momentum that is deficient only by a factor of a few. However, these cosmological simulations still fail by orders of magnitude to resolve the physics of star formation and feedback. The coupling of feedback energy with interstellar gas through shocks, the formation of galactic outflows, and the coupling of the outflow energy with halo gas are not physically represented, especially in SPH simulations with resolution at the galactic scale far too low to accurately resolve shocks. Therefore, it is important to model the collective action of multiple supernovae and study the formation of superbubbles and galactic winds in single dwarf galaxies, using high-resolution, hydrodynamic simulations, in order to compute the feedback efficiencies (e.g.\\ Mac Low \\& Ferrara 1999; hereafter MF99). Such models must also include cosmological infall of gas at high redshift, because the ram pressure of the infall may influence the evolution of the superbubbles and the galactic winds. \\subsection{Intergalactic Metals} Observations of Ly$\\alpha$-absorbing clouds reveal the presence of metals. These clouds are thought to be regions of enhanced intergalactic medium (IGM) density distant from protogalaxies (see review by Rauch 1998). To transport metals into these regions, some mechanism such as supernova-driven galactic outflows must act. The observations show that Lyman forest clouds with neutral hydrogen column density with $\\log N(HI)>14$ are metal enriched to $Z\\approx 10^{-3}-10^{-2}$ Z$_{\\odot}$ (Cowie et al.\\ 1995; Songaila \\& Cowie 1996; Ellison et al.\\ 2000), and that the metallicity remains roughly constant throughout the redshift range $1.5 < z <5.5$ (Songaila 2001). It is also interesting to note that $\\log N(HI)\\approx 14$ at $z\\approx3$ marks the transition between continuous filamentary structures and voids (e.g.\\ Zhang et al.\\ 1998). The main candidates for the polluters of the IGM are starburst dwarf galaxies. The absence of turbulent motions observed in low-density Lyman $\\alpha$ clouds at $z\\approx3$ (Rauch et al.\\ 2001b) suggests that metal enrichment was completed at very early times ($z \\gg 5$), when the physical volume of the universe was smaller, and contained the numerous dwarf galaxies predicted by the CDM model of galaxy formation (Rauch et al.\\ 2001b; Madau, Ferrara, \\& Rees 2001; Scannapieco, Ferrara, \\& Madau 2002). {\\c2 On the other hand, the median Doppler parameters are significantly larger in the Lyman $\\alpha$ forest than predicted by cosmological simulations, suggesting additional energy injection provided by late He{\\sc II} reionization or supernova-driven winds (Meiksin, Bryan, \\& Machacek 2001). In C{\\sc iv} systems with higher column densities, a substantial velocity scatter over length scales of a few hundred parsecs is observed, however, suggesting} they have been influenced by galactic feedback more recently (Rauch et al.\\ 2001a; see Rauch 2002 for summary). Galactic outflows both at very high redshift and at lower redshift $z\\sim3$ seem to play significant roles in enriching various regions of the IGM. {\\cc Cosmological simulations have suggested that the IGM at $z\\approx3$ can be metal-enriched to $Z=10^{-3}$--$10^{-2}$~Z$_{\\odot}$ by merging of or by outflows from dwarf galaxies (Gnedin 1998; Theuns et al.\\ 2002; Thacker et al.\\ 2002) without dynamically disturbing the observed, low-density, Ly$\\alpha$ clouds.} These simulations have the same problem as earlier that the feedback schemes differ based on different assumptions about the coupling of supernova energy to the surroundings, and the simulations fail badly to resolve star formation and the dynamics of bubbles. \\subsection{Starburst Wind Models} The effects of repeated supernova explosions from starbursts appear to be a central piece of physics necessary for understanding the role of stellar feedback in both galaxy formation and IGM pollution. The effects of starbursts from dwarf galaxies on the surrounding interstellar medium have been studied in the past both analytically and numerically (Mathews \\& Baker 1971; Larson 1974; Saito 1979; Dekel \\& Silk 1986; De Young \\& Heckman 1994; Silich \\& Tenorio-Tagle 1998; MF99; D'Ercole \\& Brighenti 1999). For example, MF99 studied the effects of repeated supernovae on local dwarf galaxies with gas mass $M_{g}=10^{6}$--$10^{9}$~M$_{\\odot}$ using hydrodynamic simulations, and made a parameter study of mass and metal ejection efficiencies. The main results of MF99 are that mass loss is very inefficient except in the lowest mass halos, while a substantial fraction of the hot, metal-enriched gas escapes from any of the potentials they studied. These results were confirmed by D'Ercole \\& Brighenti (1999) with a better treatment of thermal conduction. Recently, Mori, Ferrara, \\& Madau (2002) studied the same problem, but in a spherical, non-rotating, galaxy embedded in a halo with $M_h=10^{8}h^{-1}$ M$_{\\odot}$ at $z=9$. They find that, with star formation efficiency $f_* =0.1$, about 30\\% of the available supernova energy is transferred to the surrounding gas as kinetic energy. The starburst bubbles expand over $\\gg 10R_v$, where $R_v$ is the virial radius of the halo, comparable to the mean proper distance between neighboring low-mass systems. Their results support the suggestion that IGM metal enrichment occurs early. {\\cc The recent study by Wada \\& Venkatesan (2003) improves on Mori et al.\\ (2002) by precalculating a thin, dense disk with an inhomogeneous ISM in a given halo potential and by including self-gravity. They perform two three-dimensional models of a $10^8$~M$_{\\odot}$ galaxy with star formation efficiency $f_*=0.14$ and~0.014 in our terms.} We follow the work of MF99 and extend their study to high redshift. {\\cc Although the approachs taken by Mori, Ferrara, \\& Madau (2002) and Wada \\& Venkatesan (2003) are more realistic in some ways, the former neglects rotational flattening of the galaxy, and both are far more expensive computationally.} The MF99 approach enables us to do a parameter study on the feedback process by exploring a wide range of galaxy masses, formation redshifts, and star formation efficiencies. {\\cc However, note that our study is limited to axisymmetric geometry and also to single star formation sites at the centers of disks with a smooth ISM.} We perform careful, high-resolution models of the result of multiple supernova explosions in single dwarf galaxies at high redshift, using a hydrodynamic code, ZEUS-3D (Stone \\& Norman 1992; Clarke 1994), and compute the efficiency of metal and mass ejections and energy transport from the galactic potentials. We include the evolving dark matter halo potentials, and the cosmological infall of halo gas computed with a one-dimensional hydrodynamic code (Meiksin 1994). The bubbles in MF99's local galaxies were solely governed by their interaction with the interstellar medium (ISM), and so high metal ejection efficiencies were guaranteed once the bubbles blew out of the galactic disks. However, the bubbles in our high-redshift galaxies must still fight the ram pressure of the infalling halo gas after blowout before the metals, swept-up mass, and energy escape to intergalactic space. We also compute the accretion timescales for any mass swept up by the bubbles, but bound by the potential to fall back to the center. We model systems that we think likely to host a starburst. It is not certain whether first generation systems were efficient in ejecting metals and transporting energy to the IGM, because an early cosmic UV background suppreses the formation of stars inside the systems by photodissociating their only cooling agent, molecular hydrogen (Haiman, Abel, \\& Rees 2000), {\\cc or by raising the entropy floor (Oh \\& Haiman 2003)} However, later systems that collapsed due to atomic hydrogen cooling may have been more robust allowing starbursts to occur (Madau, Ferrara, \\& Rees 2001; Scannapieco, Ferrara, \\& Madau 2002; Oh \\& Haiman 2002). We choose seven of these second generation systems, with $5\\times10^{8}$~M$_{\\odot} \\le M_h \\le 5\\times10^{10}$ M$_{\\odot}$ at $3\\le z\\le 13$ with star formation efficiencies $f_{*}=0.001$, 0.01, and 0.1. Our goal is to develop a generally useful description of supernova feedback in terms of metal ejection efficiencies $\\xi_{metal}$, mass ejection efficiencies $\\xi$, and energy transport efficiencies $\\zeta$, that can be widely used in large-scale cosmological simulations for the study of galaxy formation and metal enrichment. However, it will be shown in this paper that parameterization of feedback is very difficult, and that we eventually need a higher-resolution, adaptive mesh refinement, cosmological simulation with a realistic treatment of star formation to do so as reliably as we might wish. Nonetheless, we are able to reach several key conclusions with the present set of computations. The paper is organized as follows. In \\S\\ref{sec:predict} we analytically predict the effects of ram pressure on the confinement of the starburst bubbles. In \\S\\ref{sec:numer} we discuss our numerical methods, including our models of cosmological infall (\\S\\ref{sub:halo}), starburst blowout (\\S\\ref{sub:2D}), and the tracer field we use to follow metal-enriched gas in the blowout model (\\S\\ref{sub:tracer}). In \\S\\ref{sec:disk} we describe our choices for galaxy size (\\S\\ref{sub:galaxies}), disk model (\\S\\ref{sub:disks}), and star formation feedback (\\S\\ref{sub:starbursts}). We give our results in \\S\\ref{sec:results}: metal and mass ejection, and energy transport efficiencies in \\S\\ref{sub:feedback}, the effects of external pressure on the bubble evolution in \\S\\ref{sub:press}, and the timescales for material to fall back in \\S\\ref{sub:time}. Our conclusions follow in \\S\\ref{sec:concl}. Throughout the paper, we apply $\\Lambda$CDM cosmology with $\\Omega_{0}=0.37$, $\\Omega_{\\Lambda}=0.63$, $\\Omega_{b}=0.05$, $h=0.7$, $\\sigma_8=0.8$ and a primordial composition of H:He=12:1. ", "conclusions": "\\label{sec:concl} We study feedback from dwarf starburst galaxies with virial temperatures $T_{v}>10^{4}$~K at high redshift, by numerically modeling the interaction of superbubbles driven by repeated supernova explosions with the ram pressure from cosmological infall. We compute ejection efficiencies for mass $\\xi$, metals $\\xi_{met}$, and energy $\\zeta$, after the bubbles have expanded well beyond the virial radius, and estimate the timescale for disturbed gas to fall back and become available for subsequent starbursts. We have tried to consistently choose approximations that underestimate the effects of supernova feedback, so that our computations of the efficiency of mass, metal, and energy feedback and fallback time should be reasonably strong lower limits to the actual values. We now list our conclusions. \\begin{itemize} \\item~Ejection and transport efficiencies are primarily determined by the efficiency of star formation $f_{*}$. With $f_{*}=0.1$, nearly all the metals produced in the starburst escapes ($\\xi_{met}\\approx1$), most of the disk ISM and large parts of the infalling halo are blown away ($\\xi>1$), and accelerated halo gas carries the energy of the bursts out to the IGM ($\\zeta>0.5$). With $f_*=0.01$, moderate to high feedback efficiencies are only observed in halos with $v_c\\approx30$--50 km s$^{-1}$. With $f_{*}=0.001$, none of the metals, the mass, nor the energy can escape the potentials. \\item~The ram pressure of infalling halo gas can suppress the growth of starburst bubbles and thus prevent the metal, mass, and energy feedback, unless $f_{*} \\ga 0.1$. The energy available to drive the bubble depends on both $f_*$ and on the fraction of energy in the hot gas actually available to drive the bubble $\\nu$. The minimum halo mass for ram pressure confinement with a given $\\nu f_{*}$ can be specified by a cut-off in virial temperature $T^{*}_{v}$ or circular velocity $v^{*}_c$ at any redshift (Eq.~\\ref{tconf}). Our prediction is consistent with the values of $\\nu$, $\\xi_{met}$, $\\xi$, and $\\zeta$ found in our simulations for a given $f_*$. With $\\nu f_{*}=1\\times10^{-3}$, we can only expect high feedback efficiencies in halos with $T_{v}\\la T^{*}_{v}\\approx 1\\times10^{5}$ K and $v_c\\la v^{*}_c\\approx 52$ km s$^{-1}$, above the limit for hydrogen line cooling, and below the Jeans mass in the reheated IGM. \\item~We find that small dwarf galaxies with $v_c\\approx30$--50 km s$^{-1}$ at $z\\la8$ are efficient in enriching the IGM with metals, yielding $\\bar{Z}\\ga10^{-3}$ Z$_{\\odot}$ with $f_*=0.01$. The early metal enrichment may be able to explain the absence of kinematic disturbance observed in low-density Lyman $\\alpha$ clouds. Larger dwarf galaxies with $v_c\\approx100$ km s$^{-1}$ at $z=3$--5 might enrich the IGM more effectively, yielding $\\bar{Z}\\approx10^{-2}$ Z$_{\\odot}$, because efficient cooling expected in such halos may enable $f_{*}\\gg0.01$. The late metal enrichment is consistent with high metallicities and large turbulent motions observed in high-density Lyman $\\alpha$ clouds. These galaxies form out of 1--2$\\sigma$ peaks in our $\\Lambda$CDM Gaussian density perturbation. \\item~We expect mass ejection efficiencies $\\xi\\approx0.3$--1 and energy transport efficiencies $\\zeta\\approx0.1$-0.3 from the dwarf galaxies with $v_c\\approx30$--50 km s$^{-1}$ with $f_*=0.01$. However, most of the ISM remains bound unless $f_{*}=0.1$. Instead, a large amount of infalling gas is turned around by the galactic wind and escapes the potential. This outflow may carry kinetic energy into the IGM, broadening Ly$\\alpha$ absorption systems, and to other nearby halos to prevent gravitational collapse of the gas in them. We estimate that the ensemble of outflows may provide turbulent support with typical $v_{rms}\\approx20$--40 km s$^{-1}$. Further study of this multiscale problem is required to determine whether this mechanism of suppressing dwarf galaxy formation can solve the overcooling and angular momentum problems. \\item~The timescale for swept-up mass to fall back to the center of a halo is short compared to the Hubble time, $\\tau_f \\la 0.4 t_H(z)$ if $f_{*}\\la0.01$. We expect star formation efficiencies for subsequent starbursts to be larger than those of the initial starbursts that we studied, because ejected metals bound in the potentials will enhance gas cooling. Once gas in a halo cools to form a disk-like object, it is difficult to blow away the disk gas. \\end{itemize} We must mention the caveats in our study: Our models assume an instantaneous starburst at the center of a galactic disk. This is obviously an oversimplification, and such instantaneous starbursts, especially with $f_{*}\\gg0.01$ are probably not realistic. We instead expect multiple OB associations or star clusters to form over a few hundred pc across the disk (e.g. Vacca 1996; Martin 1998). Each cluster will likely behave like our models however, so the net effect will be greater than we compute, consistent with our attempt to present lower limits to the effects of stellar feedback. Given the dominant influence of the star formation efficiency on the importance of stellar feedback, a self-consistent, physically motivated modeling of star formation will be required to further our understanding. We entirely neglected magnetic fields in this study. Magnetic fields can inhibit the formation of cold, dense shells, and suppress the fragmentation of the shells, perhaps reducing or preventing the dissipation of the bubble energy through the fragmented shells. Then the fraction of energy preserved in the bubbles $\\nu$ might be larger, giving higher feedback efficiencies in the dwarf galaxies that we modeled. On the other hand, magnetic pressure and tension might act to help confine the expanding bubbles (e.g.\\ Tomisaka 1998), depending on the strength of the field. Our models are not three-dimensional. This has consequences for two main reasons. First, we compute the Rayleigh-Taylor instability in fragmenting shells during blowout in two dimensions. This has the effect of increasing the size of the fragments (Mac Low et al.\\ 1989) because only ring-like modes of the instability can form. In three dimensions, the instability takes on a spike and bubble form, with the spikes representing the dense fragments. This does not change the gross topology, since even in two dimensions the shell still overturns and the wind can escape. However, a detailed study of the fragment properties cannot be done. As we restrict ourselves to integrated quantities, however, this probably does not strongly affect our results. Certainly the two models computed by Wada \\& Venkatesan (2003) in three dimensions appear to grossly agree with our results: a star formation efficiency of $f_* = 0.14$ in a galaxy with $10^7$~M$_{\\odot}$ of gas leads to metal escape, while with $f_* = 0.014$ the metals remain confined. Second, our models assume spherically symmetric gas infall, and the presence of a well-formed, rotationally supported disk. It is also an oversimplification to compute metal, mass, and energy feedback efficiencies when the bubbles reach $R_s=2R_v$. Realistically, a halo is embedded in the complicated weblike structures predicted by three-dimensional hydrodynamic simulations (e.g.\\ Miralda-Escude et al.\\ 1996; Zhang et al.\\ 1998). Even inside its virial radius, gas accretion occurs along filaments, and the accretion shock remains highly aspherical even after the protogalaxy has collapsed to less than a few times the virial radius, (Abel et al.\\ 1998; 2000; 2002). In addition, a rotationally supported disk may not form before a significant amount of star formation occurs. We note, however, that a more dynamical, filamentary background will create funnels through which superbubbles and galactic winds can freely expand, potentially increasing the efficiency of metal ejection and energy transport. Although these caveats identify important points requiring further investigation, none of them appear to invalidate our results; nor do they suggest that our lower limits are drastic underestimates of the actual effects of stellar feedback." }, "0405/astro-ph0405198.txt": { "abstract": "{We present and discuss new determinations of metallicity, rotation, age, kinematics, and Galactic orbits for a complete, magnitude-limited, and kinematically unbiased sample of 16,682 nearby F and G dwarf stars. Our $\\sim$63,000 new, accurate radial-velocity observations for nearly 13,500 stars allow identification of most of the binary stars in the sample and, together with published $uvby\\beta$ photometry, Hipparcos parallaxes, Tycho-2 proper motions, and a few earlier radial velocities, complete the kinematic information for 14,139 stars. These high-quality velocity data are supplemented by effective temperatures and metallicities newly derived from recent and/or revised calibrations. The remaining stars either lack Hipparcos data or have fast rotation. \\\\ A major effort has been devoted to the determination of new isochrone ages for all stars for which this is possible. Particular attention has been given to a realistic treatment of statistical biases and error estimates, as standard techniques tend to underestimate these effects and introduce spurious features in the age distributions. Our ages agree well with those by Edvardsson et al. (\\cite{edv93}), despite several astrophysical and computational improvements since then. We demonstrate, however, how strong observational and theoretical biases cause the distribution of the {\\it observed} ages to be very different from that of the {\\it true} age distribution of the sample.\\\\ Among the many basic relations of the Galactic disk that can be reinvestigated from the data presented here, we revisit the metallicity distribution of the G dwarfs and the age-metallicity, age-velocity, and metallicity-velocity relations of the Solar neighbourhood. Our first results confirm the lack of metal-poor G dwarfs relative to closed-box model predictions (the ``G dwarf problem''), the existence of radial metallicity gradients in the disk, the small change in mean metallicity of the thin disk since its formation and the substantial scatter in metallicity at all ages, and the continuing kinematic heating of the thin disk with an efficiency consistent with that expected for a combination of spiral arms and giant molecular clouds. Distinct features in the distribution of the $V$ component of the space motion are extended in age and metallicity, corresponding to the effects of stochastic spiral waves rather than classical moving groups, and may complicate the identification of thick-disk stars from kinematic criteria. More advanced analyses of this rich material will require careful simulations of the selection criteria for the sample and the distribution of observational errors. ", "introduction": " ", "conclusions": "The Galaxy is a far more complicated and interesting subject than ever before. The present work should lay the foundation for learning more about it. Looking ahead a decade from now, the ESA cornerstone mission GAIA (Perryman et al. \\cite{gaia01}) will provide the next quantum leap in our knowledge of the Galaxy. Obtaining the complementary photometry and radial velocities needed to fully exploit the astrometric data from Hipparcos was left to such independent ground-based efforts as the present programme (see also Udry et al. \\cite{udryetal97}). In contrast, GAIA will obtain these observations with the same satellite payload as the astrometric data, although with much larger radial-velocity errors than achieved here (cf. Fig. \\ref{vrstat}). The present material should remain useful until the results from GAIA appear, not only for studying our Galaxy, but also in the efforts to optimise the observing and data reduction strategies for GAIA and for such precursor programmes as the RAVE survey (Steinmetz \\cite{steinmetz})." }, "0405/astro-ph0405082_arXiv.txt": { "abstract": "{ We present the discovery of a binary companion to the pulsating sdB Feige~48. Using HST/STIS and archival FUSE spectra, we measure a period of 0.376$\\pm$0.003\\,d and a velocity semi-amplitude of 28.0$\\pm$0.2\\,km\\,s$^{-1}$. This implies that the companion star must either be of very low mass, or the orbit is at high inclination. Combining \\emph{2MASS} fluxes, the lack of a reflection effect, results from asteroseismology and a measurement of the rotation velocity of Feige~48, we show that the orbital inclination must be $\\le$11.4$^\\circ$ and that the unseen companion is a white dwarf with mass $\\ge$0.46\\,$M_\\odot$. The implications of this discovery, and of binarity amongst sdB pulsators, is then discussed in the context of recent theoretical work on sdB formation. In particular we suggest that radial velocity studies focus on sdB pulsators with no known companion, and that asteroseismological studies of sdBs investigate a larger mass range than previously considered in order to test formation models. ", "introduction": "\\label{sec:intro} In recent years, interest in hot subdwarf B (sdB) stars has increased dramatically. This is in large part due to the discovery of rapid oscillations in some sdBs, but also because improving instrumentation has allowed more detailed studies of their spectral properties. It is now generally accepted that sdBs are associated with stars on the extreme Horizontal Branch (EHB) \\citep[e.g.,][]{Heber86, SBK94}. This means they are low-mass stars ($\\sim$0.5\\,$M_\\odot$) that burn helium in their cores and have hydrogen envelopes too thin ($<$0.02\\,$M_\\odot$) to sustain nuclear burning. Stellar evolution calculations by \\citet{DRO93} predict that because of their thin envelopes, sdBs will evolve directly to the white dwarf cooling track, bypassing the Asymptotic Giant Branch. While the next stages of sdB evolution appear to be known, the question of sdB formation remains unanswered. There have been three basic scenarios proposed: binary interaction involving Roche lobe overflow \\citep{MNG76}, single star evolution with strong mass loss near the tip of the Red Giant Branch \\citep{DCru96b}, and the merger of two helium white dwarfs \\citep{IT1986}. Significant progess has been made with the recent work of \\citet{HanI,HanII}, who studied three possible binary formation channels: common envelope ejection, stable Roche lobe overflow and the merger of two helium white dwarfs. Using binary population synthesis, they created 12 simulations and found one of them could satifactorily reproduce the observed characteristics of sdBs ($T_{\\mathrm{eff}}$, $\\log~g$, binary fraction, orbital period distribution, etc) when compared to radial velocity studies by \\citet{MHMN01} and \\citet{M-RMM03}. \\citet{RKL04} found a binary fraction of 39\\%, lower than predicted, while \\citet{Lisker2004}, in the largest homogenous study so far, compared their data to the simulations of \\citet{HanII} and found a different best-fit model. The oscillations detected in some sdBs offer a way to test these formation models. Importantly, pulsations are seen in sdBs with short- and long-period binary companions, as well as in apparently single sdBs. The properties of the V361\\,Hya stars (the short-period, $p$-mode sdB pulsators, originally called EC\\,14026 stars) are reviewed by \\citet{Kilkenny2002}, while \\citet{GFR03} presents the detection of long-period, likely $g$-mode pulsations in somewhat cooler sdBs (currently known as the PG\\,1716+426 or ``Betsy'' stars). Pulsations in the V361\\,Hya star, Feige~48, were first detected by \\citet{ECpaperXI}. It is one of the coolest sdB pulsators in its class, with $T_{\\mathrm{eff}}$=29\\,500\\,K and $\\log g$=5.50 \\citep{HRW00}, and was discovered to be pulsating with five periods between 344 and 379\\,s. The oscillation amplitudes show dramatic variation from night to night, which the authors suggested are due to either changes in pulsation power, the beating of unresolved modes, or a combination of the two. In a comprehensive follow-up study, \\citet{RKZ04} observed Feige~48 over 5 years, and found that the periods measured by \\citet{ECpaperXI} were resolved and that the amplitudes vary by at least 30\\%. A tentative model match to four out of five detected periods, suggests that Feige~48 rotates with a period around 10\\,h. The $T_{\\mathrm{eff}}$ and $\\log g$ of this model were chosen to be consistent with those derived from spectroscopy. Long-term monitoring by \\citeauthor{RKZ04} also provided an opportunity to measure the phase stability of the pulsations. In particular, they were able to place constraints on the period and mass of a companion to quite high precision, limiting any stellar-mass companion to have a period shorter than 3 days. In another recent study, \\citet{CBF03} found a very good model match to all nine periods detected by high S/N photometry using the CFHT. They derived a rotational period of 9.58\\,h, in good agreement with \\citeauthor{RKZ04} In this paper we present the detection of binary motion in Feige~48, an sdB that was generally assumed to be single. We discuss the nature of the companion, and implications of this detection for current theories of sdB formation and pulsation. ", "conclusions": "\\label{sec:disc} We have detected binary motion in the V361\\,Hya star Feige~48 with a period of 0.376\\,d and a velocity semi-amplitude of 28.0\\,km\\,s$^{-1}$. Combining multicolour photometry, the asteroseismological determination of the rotation period, and an upper limit of the equatorial rotation velocity, we conclude that the orbital inclination $i\\le$11.4$^\\circ$. This implies that the minimum mass of the companion is 0.46\\,$M_\\odot$; combining this with the lack of a reflection effect suggests the star must be a white dwarf. Feige 48 is one of only two pulsating sdBs with a known white dwarf companion. Unlike KPD\\,1930+2752, the other system, it has a relatively simple pulsation spectrum and no (or very little) ellipsoidal distortion. This has allowed \\citet{MR_PhD02}, \\citet{RKZ04} and \\citet{CBF03} to match pulsation frequencies and constrain sdB evolution models. However, these studies assumed that Feige 48 evolved according to canonical EHB theory. While this is probably accurate since the star has undergone two phases of common envelope evolution, the conclusions of \\citet{HanII} suggest that investigating a wider mass range than previously may be warranted, with masses as low as $\\sim$0.36\\,$M_\\odot$. In the case of single sdBs, which should be formed through the merger channel, masses in the range 0.4-0.7\\,$M_\\odot$ should be examined. The sdBs in long-period orbits with main sequence companions (formed through the Roche lobe overflow channel) may fall in the 0.3-0.5\\,$M_\\odot$ range, and possibly even higher, although \\citeauthor{HanII} find this to be unlikely. Enlarging the mass range investigated in this way may provide another useful test of the proposed formation models for sdBs. These tests can only be carried out once the binary nature of the pulsating sdBs is known. Up to now only seven V361\\,Hya stars are confirmed binaries; four have F- or G-type main sequence companions (V361\\,Hya itself, PB\\,8783, EC\\,20117-4014 and EC\\,10228-0905), two have white dwarf companions (KPD\\,1930+2752 and Feige~48), and one is in an eclipsing binary system with an M5 main sequence star (PG\\,1336--018). That leaves around 25 pulsators with unknown status. The fraction of sdBs is binary systems is still under debate, however if we consider the best-fit simulation of \\citet{HanII}, we find that $\\sim$41\\% or 10-11 of these stars should have either late-type main sequence or white dwarf companions. \\citet{Lisker2004} compared their observations of 52 sdB stars with the 12 simulation sets constructed by \\citet{HanII} using a simple statistical test. The simulation set they found to fit their observations differs from that found qualitatively by \\citeauthor{HanII} and suggests that up to $\\sim$53\\% or $\\sim$13 of the pulsating sdBs should have an unseen companion. \\emph{High precision spectra are needed to observe short period systems such as Feige~48, so a survey with high spectral and temporal resolution is required.}" }, "0405/astro-ph0405561_arXiv.txt": { "abstract": "The role of stars and starbursts in AGN has been a recurring issue for nearly as long as AGN have been recognized as hosts of interesting phenomena. The heated ``starburst {\\it versus} monster'' controversy of the 80's and 90's was gradually replaced by ``starburst {\\it plus} monster'' studies, as observational work in the past decade has firmly established that accretion onto a super-massive black-hole and star-formation coexist in many galactic nuclei. Whereas the physical link between starbursts and AGN remains unclear, there remains no doubt that starbursts affect a number of properties traditionally associated to the AGN alone, such as the so called ``featureless continuum'', emission line ratios and luminosities. This contribution glosses over some of the techniques used to diagnose stellar populations in AGN, focusing on recent results and how this type of work can lead us well beyond what became known as the starburst-AGN connection. ", "introduction": "The subject of stellar populations in AGN has gone through a history of love and hate over the past 3 decades. Once upon the time, stars were essentially seen as that unavoidable junk which pollutes the optical spectra of AGN, particularly type 2s (Seyfert 2s, LINERs and their relatives). Accordingly, the methodology to deal with starlight in those days was philosophically the same used to deal with sky features or cosmic rays: {\\it Get rid of it!} This was achieved by modeling the spectrum as a combination of an elliptical galaxy template plus a non-stellar featureless continuum (FC), yielding the ``pure-AGN'' spectrum as the residual (Koski 1978). Since this approach postulates that stars in AGN are all old and boring, it is not surprising that few AGNauts cared about stellar populations at all, so meetings like these would not have been possible at the time. This ``Get Rid of It'' era lasted up to the mid-80's. In the late 80's and 90's things changed radically from a ``stars-have-nothing-to-do-with-AGN'' to the ``stars = AGN'' idea put forward by R. Terlevich and collaborators, whose starburst model for AGN replaced super-massive black-holes (SMBHs) by young stars and their remnants as the power-engine of active galaxies. Those of use which were around at the time remember (some with nostalgia) the fierce debates in meetings across the globe (eg, Taipei 1992, Puebla 1996). The conflicting observational evidence back then fueled the controversy. For instance, while the discovery of Seyfert 1-like SNe supported the model, rapid X-ray variability and relativistic Fe K$\\alpha$ line-profiles were better understood in the framework of the black-hole paradigm. The starburst $\\times$ monster battle gradually disappeared from the headlines as the existence of SMBHs became conclusively established in the past decade. The interest in stellar populations in AGN, however, survives to the present day, and for a very good reason: Evidence of the presence of young and intermediate age population around AGN is now as solid as that for the existence of SMBHs! Young ($t < 10^7$ yr) and ``mature'' ($10^{8 {\\rm -} 9}$ yr) starbursts have been found all across the AGN family: quasars (Brotherton \\etal 1999), radio galaxies (Wills \\etal 2002), Seyfert 2s (Heckman \\etal 1997) and LLAGN (Cid Fernandes \\etal 2004; Gonz\\'alez-Delgado \\etal 2004). This very volume reports new developments on this so called ``starburst-AGN connection'', as in the contributions by Rafaela Morganti, Wil van Breugel, Rosa Gonz\\'alez Delgado and others. Indeed, the mere fact that the word ``stars'' figures alongside ``black-holes'' in the title of an IAU symposium is a testimony of this new reality. Having said that, I must point out that we are still a long way away from understanding the physical link between star-formation and SMBHs in AGN. In fact, except for a few cases (see van Breugel's contribution) we are not even sure such a causal link exists at all, as the coexistence of these two phenomena could simply reflect the fact that both live on the same gas-based diet, a trivial possibility one must always bear in mind. Notwithstanding this obligatory warning, we have learned a great deal from such studies. Instead of attempting a thorough and fair review of the progress in the field, which would be impossible due to lack of space and mainly talent, these few pages highlight some recent results which illustrate how the careful modeling of stellar populations in AGN provides valuable tools both for researchers interested in studying starburst-AGN connections and those more interested in getting rid of the starlight to inspect the AGN itself. The work cited below was carried out with several friends, including R. Gonz\\'alez Delgado, H. Schmitt, T. Heckman, L. Martins, Q. Gu, J. Melnick, the Terlevichs, D. Kunth, my students and, last but not least, our chair-woman T. Storchi-Bergmann, who taught me a lot and worked very hard to organize this great meeting. ", "conclusions": "" }, "0405/hep-ph0405218_arXiv.txt": { "abstract": "{\\small We propose a new method for identifying new physics imprints on present observational data in cosmology whereby signatures of string theory are clearly distinguished from imprints of possible features on the inflaton potential. Our method relies on the cross-correlations spectra of cosmic shear from large scale structure (LSS) with the CMB temperature anisotropies and E-mode polarization, by using the following properties: inflationary cosmology provides only one source term for all CMB spectra and LSS which highly constrains any deviations from the standard predictions; string theory can add new non-inflationary channels to the source of perturbations as well as modify clustering properties at large scales. Discrepancies in the source terms of correlations and clustering properties provide the evidence for new physics. Models of single-field inflation with a feature are disfavored even with present data. Upcoming WMAP results and future data from weak lensing of LSS will further improve our ability to probe new physics in this manner and could open the first direct window to string theory.} ", "introduction": "There is an emerging picture of the universe obtained from the tremendous progress achieved by precision cosmology. According to this picture we live in a 'weird' universe. A universe which, at present redshifts, $z \\simeq 0-1$, and energy scales $H_0 \\simeq 10^{-33}eV$, is: dominated by a mysterious component of energy, coined dark energy, driving it into an accelerated expansion \\cite{de};and which, contrary to our theoretical expectations based on the inflationary paradigm, has large angle CMB perturbations suppressed \\cite{wmap}. The later puzzle is not confirmed yet due to limitations from cosmic variance \\cite{cosmicvar}, and the need for a better understanding of the systematic and foreground effects. Hence a lot of work done in investigating this effect \\cite{cross, kesden}, by using complementary data to minimize statistical errors, points out that the suppression of CMB power for large angles $\\theta$ persists. In what follows, we take the WMAP CMB measurements \\cite{wmap} at large scales as true and consider the power suppression at low multipoles $l$ to be a real physical effect rather than a statistical fluke. Dark energy domination and CMB power suppression of perturbations at large angles, both occur, around the same redshift and energy scale $H_0 \\simeq 10^{-33}$ eV. They both raise two disturbing questions: Why is their magnitude so small, (the tuning problem), and, why are they occurring now, (the coincidences problem). A consistent theoretical model should simultaneously address both questions, the tuning and the cosmic coincidences, for these observed phenomena. Besides, it should also address why both coincidences occur at the {\\it same} energy scale, our present Hubble radius $H_0$, \\cite{ale}. It is tempting to speculate that too many coincidences in the present universe may indicate the emergence of a new scale in physics of the order our Hubble scale $H_0 \\simeq 10^{-33}$ eV. No such model exists yet. Perhaps because an understanding of the underlying fundamental theory is required in order to explain the current challenges that theoretical cosmology is facing. As yet string theory is the leading candidate for new physics although it has still not provided us with satisfactory answers to the above issues. A conservative approach would be to question whether we need new physics at all for addressing the current problems. Distinct observational signatures would be the best way for resolving these doubts and for providing direct evidence of new physics if it exists. Therefore, it is of crucial importance to attempt to distinguish whether what is giving rise to the observed cosmic coincidences are features of inflaton \\cite{easther,linde} or imprints of new physics \\cite{efstathiou,katie,costa}. This question is the scope of our work here, namely: Can we uniquely identify and discriminate new physics signatures from inflaton features in the presently available observables in the sky? A positive answer to the above question in favor of new physics, possibly string theory as its leading candidate, would provide the first direct evidence for its existence. This is a difficult task to investigate. Any signatures expected from new physics, e.g. anomalies on the CMB spectra, are very small\\cite{tp}. Furthermore, even if they are within the detection limits, it is notoriously difficult to identify whether it is some feature of the unknown inflaton potential\\cite{easther} that give rise to the anomalies or whether the anomalies are rooted in new (string) physics. Besides, the origin of dark energy and its relation to the power suppression of low multipoles around redshift $z \\simeq 0$ is not yet understood. In this situation, we need to identify an imprint which arises uniquely from one type of theory but is generically negligible in the other. We demonstrate in the next section that finding such a unique handle on the observational signatures is possible and then use this handle to discriminate among various models with respect to their imprints on the combined data from CMB spectra, LSS and dark energy. The autocorrelation spectra do not yield any new information on the origin of the observed anomalies since they are sourced by the same term \\cite{tp}. In the standard theory of inflation and physics of the early universe, the temperature perturbations given by the Sachs-Wolfe effect are \\begin{equation} \\label{temp1} \\Theta({\\bf n}) = \\frac{1}{3} \\Phi(r_0,z_0) - 2\\int_0^{r_0} \\frac{d\\Phi}{dr}(r,z(r))dr \\,\\, \\end{equation} where $\\Phi(r,z)$ is the gravitational background potential and $r,z$ are the physical comoving distance and redshift respectively. The index $0$ denotes the value at last scattering surface. The gravitational potential $\\Phi(r,z)$ is related to the growth factor of structure $G(z)$ through $\\Phi(r,z) = (1+z)G(z)\\Phi(r,0)$.Denoting the three dimensional matter power spectrum by $P(k)$ then the power spectrum for the CMB temperature anisotropies in a flat universe is \\begin{equation} \\label{temp2} C_l^{TT} \\simeq \\int \\frac{dk}{k^2} P(k) [\\Theta_l (k)]^2 \\, \\end{equation} where $\\Theta_l (k)$ is the Fourier transform of $\\Theta({\\bf n})$. An expression similar to Eqn.~\\ref{temp2} gives the spectrum of the curl free E-mode of polarization of CMB photons, denoted by $C_l^{EE}$ \\cite{ma}. But the same gravitational potential $\\Phi(r,z)$ that produces the temperature anisotropies $C_l^{TT}$ also sources the polarization spectra namely the $E, B$ -modes and gravitational lensing of LSS. In fact gravitational lensing can be a powerful tool for mapping the background potential for LSS. The projected spectrum is described by the following potential \\begin{equation} \\label{lens} L({\\bf n}) = -2\\int_0^{r_0} dr \\frac{r_0 -r}{r r_o}\\Phi(r,z(r)) \\, \\end{equation} and the angular power spectrum of the lensing potential is \\begin{equation} \\label{lens2} C_l^{LL} \\simeq \\int \\frac{dk}{k^2} P(k) [L_l (k)]^2 \\, \\end{equation} where similarly $L_l (k)$ is the Fourier transform of $L({\\bf n})$. Notice however that although both spectra depend on the same source in the conventional theory, their integral dependence on $\\Phi(r,z)$, in the expressions (\\ref{temp1}) and (\\ref{lens}) are sufficiently different. This fact can be exploited by using the cross correlation spectrum between $T,L$ in identifying disreptancies with the expectations of the standard concordance cosmology. \\begin{equation} \\label{crossshear} C_l^{TL} \\simeq \\int \\frac{dk}{k^2} P(k)[ \\Theta_l (k) L_l (k) ] \\, \\end{equation} A similar expression to Eqn.~(\\ref{crossshear}) gives the cross-correlation spectrum $C_l^{TE}$ of T with E-mode. The role of dark energy on the low multipoles $l$ is to enhance their power through the Integrated Sachs-Wolfe (ISW) effect. The observed suppression of power may indicate that clustering properties at large scales are different from the predictions of the standard theory and such that they compete with the ISW effect. The correlations reveal whether either of the two spectra deviate, due to new physics, from the standard dependence on $\\Phi$ expected by the equations above. Deviations may arise from contributions from noninflationary channels or a fundamental string scale imprinted on $C_l^{TL,TE}$ that may signal a breakdown of the conventional theory. Below we use the combination of both correlation spectra $C_l^{TE}$ and $C_l^{TL}$ as our handle in identifying the origin of the imprints on observations while avoiding some of the degeneracies among cosmic parameters. Generically, the inflaton field, (in single-field inflation models), is very weakly coupled to other sectors due to the slow roll conditions and adiabaticity requirement. Meanwhile there is no reason why the coupling of the stringy moduli fields to the matter sector should be suppressed. Therefore any features on the inflaton potential and their signature to LSS and clustering properties at large scales would be very different from the effects from variations of moduli couplings carried over to the matter sector. These moduli couplings could have a large impact onto the clustering properties of large scale structure and the polarization spectrum. For this reason discrepancies in the cross correlations would reveal complimentary information in identifying the origin of the observed signatures. In Sect.2 we provide the theoretical framework for a generic classification of the possible effects of string and brane-world models on the low-energy universe. There we describe how cross-correlation spectra provide the extra information needed in order to discriminate between a standard inflationary model with a feature on its potential, from the various string-inspired classes of cosmology beyond the concordance model. We focus only on the string models from literature that can give rise to an accelerated expansion of the universe at late time, a criteria imposed by observations, which takes account of the ISW effect. Data analysis of $C_l^{TT, EE, TE, TL}$ spectra for the string effects in cosmology such as variations of the string coupling constants, dispersion relation for the matter fields arising from their coupling to moduli fields, changes in the strength of gravity, large scale clustering properties and variations of the background gravitational potential $\\phi$ due to the higher dimensional nature of gravity, is presented in Sect.3. There we show how the $C_l^{TE}$ spectrum breaks some of the degeneracy of the cosmological parameters while the $C_l^{TL}$ spectrum carries the unique signature for new physics. Results obtained in Sect.3. are discussed in Sect.4. ", "conclusions": "" }, "0405/astro-ph0405031_arXiv.txt": { "abstract": "We present $5\\psec2\\times 2\\psec6$ resolution interferometry of CO \\jon\\ emission from the starburst galaxy NGC 253. The high spatial resolution of these new data, in combination with recent high resolution maps of \\ico, HCN and near-infrared emission, allow us for the first time to link unambiguously the gas properties in the central starburst of NGC 253 with its bar dynamics. We confirm that the star formation results from bar-driven gas flows as seen in ``twin peaks'' galaxies. Two distinct kinematic features are evident from the CO map and position-velocity diagram: a group of clouds rotating as a solid body about the kinematic center of the galaxy, and a more extended gas component associated with the near-infrared bar. We model the line intensities of CO, HCN and \\ico\\ to infer the physical conditions of the gas in the nucleus of NGC 253. The results indicate increased volume densities around the radio nucleus in a twin-peaks morphology. Compared with the CO kinematics, the gas densities appear highest near the radius of a likely inner Linblad resonance, and slightly lead the bar minor axis. This result is similar to observations of the face-on, twin-peaks galaxy NGC 6951, and is consistent with models of starburst generation due to gas inflow along a bar. ", "introduction": "Bar-driven gas flows have been proposed to explain the high star formation rates in the nuclei of some starburst galaxies \\citep*{twinpeaks, knapen, kohno, sheth}. Star formation is predicted to be enhanced in regions with little shear and/or weak or no shocks, i.e., where the velocity gradients are small. These regions are predicted to be found, and in some cases are observed, at the bar ends, along the leading edges of the bar, and at locations within a nuclear ring or spiral. Molecular gas concentrations often appear in these areas, in particular at so-called ``twin peaks'' near the contact points between the dust lanes of the bar and the nuclear ring \\citep{twinpeaks}. High resolution images of the molecular gas in the archetypal starburst galaxy NGC 253 suggest that a similar ring of clouds may exist around its nucleus. Specifically, the position-velocity diagram (PVD) shows a hole near the center \\citep{n253.hcop, n253.cs, n253.ihcop}, though this feature does not unambiguously indicate the presence of a ring \\citep{sakamoto}. From CS observations, \\citet{n253.cs} suggest that much of the dense, star-forming gas in the nucleus of NGC 253 lies in a ring associated with so-called $x_2$ orbits, which are small, elongated orbits perpendicular to the bar \\citep{bt}. Similarly, \\citet{n253.ihcop} conclude from the PVDs of \\ihcop\\ and SiO emission in NGC 253 that the dense gas also follows these orbits. The ratio of HCN and CO intensities, an indicator of density, is also elevated 5--$10''$ on either side of the nucleus along the major axis \\citep*{n253.hcn}. Indeed, dense gas in the face-on galaxy NGC 6951 appears to follow the $x_2$ orbits closely \\citep{kohno}, but the low declination and high inclination of NGC 253 \\citep[$87\\pdeg5$,][]{pence} have hindered similar analysis. Previous high resolution maps of gas emission from NGC 253 either did not trace the bulk of the molecular gas, i. e., they were most sensitive to dense, shocked or ionized gas \\citep{n253.hcn, n253.h92a, n253.cs, n253.ihcop}, or they lacked the spatial resolution to isolate individual cloud complexes and their kinematics \\citep*{canzian, n253.co21}. Further, since the first CO interferometer map of NGC 253 by \\citet{canzian}, new images have been made of the stellar and dust distributions \\citep{2mass} and optically thin gas traced by \\ico\\ \\citep{n253.ico}. \\citet{das} have also analyzed the CO velocity field at high resolution, and with the Submillimeter Array now operational, interferometric maps of dense and warm gas traced by CO \\jth\\ and other lines are certainly imminent. Given these new and anticipated data, it is appropriate to reexamine the properties of the molecular gas in NGC 253 but on the spatial scales of individual giant molecular clouds ($\\lsim 40$ pc). To this end, we imaged the CO \\jon\\ emission from NGC 253 with the Nobeyama Millemeter Array. To place any new work in the proper context will require this relied-upon tracer of the bulk of the star-forming gas in galaxies. We compare these data to high resolution HCN and \\ico\\ maps \\citep{n253.hcn, n253.ico} to infer the properties of the molecular clouds and compare them to the kinematics and dust distribution to test how the bar in NGC 253 may drive the central starburst. ", "conclusions": "We present the CO distribution and kinematics in the starburst nucleus of NGC 253. The molecular gas is concentrated in several complexes which are associated with both the near-infrared bar and likely ``twin peaks'' around the radio nucleus. The gas densities, temperatures and column densities are highest in these presumed twin peaks. The CO kinematics indicate that the density peaks may also coincide with an inner ILR just beyond the region of solid body rotation around the nucleus. Even without an ILR, the velocity gradient decreases near the peaks as well, which could also explain the star formation there. Analysis of the rotation curve and expected gas motions in a bar potential, compared with the locations of enhanced gas density, lead us to conclude that the stellar bar drives the star formation in the starburst nucleus of NGC 253. This conclusion is supported by the likely presence of at least one ILR, gas motions consistent with $x_1$ and $x_2$ orbits, leading dust ridges along the infrared bar, and an apparent twin peaks morphology." }, "0405/astro-ph0405494_arXiv.txt": { "abstract": "We present an extension of the harmonic-space maximum-entropy component separation method (MEM) for multi-frequency CMB observations that allows one to perform the separation with more plausible assumptions about the receiver noise and foreground astrophysical components. Component separation is considered in the presence of spatially-varying noise variance and spectral properties of the foreground components. It is shown that, if not taken properly into account, the presence of spatially-varying foreground spectra, in particular, can severely reduce the accuracy of the component separation. Nevertheless, by extending the basic method to accommodate such behaviour and the presence of anisotropic noise, we find that the accuracy of the component separation can be improved to a level comparable with previous investigations in which these effects were not present. ", "introduction": "\\label{intro} An important stage in the reduction of CMB anisotropy data is the separation of the astrophysical and cosmological components. Several techniques have been suggested, including blind (Baccigalupi \\et 2000, Maino \\et 2002) and non-blind (Hobson \\et 1998, Bouchet and Gispert 1999, Stolyarov \\et 2002) approaches. Non-blind methods, such as the maximum-entropy method (MEM) or Wiener filtering, allow one to use all available prior information about the components in the separation process. A detailed description of the harmonic-space MEM approach for flat patches of the sky was described by Hobson \\et (1998), and was extended later to the sphere by Stolyarov \\et (2002; hereinafter S02). Accounting for the presence of point sources was discussed by Hobson \\et (1999) for the flat patches and point source detection on the full sky maps using Spherical Mexican Hat Wavelets (SMHW) was analyzed by Vielva \\et (2003). A joint technique using both SMHW and MEM for the flat case was investigated by Vielva \\et (2001), and for the spherical case it will be described in a forthcoming paper by Stolyarov \\et (in preparation). The method has also been used to construct simulated all-sky catalogues of the thermal Sunyaev-Zel'dovich effect in galaxy clusters (Geisb\\\"usch, Kneissl and Hobson, in preparation). The separation tests described in previous articles were performed making some simplifying assumptions. In particular, the receiver noise was assumed uncorrelated and statistically homogeneous over the sky, which is a reasonable approximation for some scanning strategies. Another simplification concerned the foreground spectral behaviour. It was assumed that spectral parameters, such as the synchrotron spectral index and dust emission properties, were spatially-invariant. In the analysis of real data, however, one cannot simply ignore these effects, and it is very important to investigate their influence on the component separation results. For real observations, each point in the sky is observed a different number of times depending on the scanning strategy of the instrument. In the case of simple scanning strategies, such as constant latitude scans, the spin axis stays close to the plane of the ecliptic. In this situation, pixels near the ecliptic pole are observed several times more often than those near the ecliptic plane. This uneven coverage leads to marked differences in the noise rms per pixel across the sky. We also note that instrinsic gain fluctuations in the instrument can also contribute to the variable noise rms. Component separation using incomplete maps containing cuts (e.g. along the Galactic plane) can be considered as an extreme case of varying noise rms. One approach is to assume that the noise rms for pixels in the cut is formally infinite (or, equivalently, that they have zero statistical weight). We show below that this method does indeed allow the separation method to cope straightforwardly with cuts. Moreover, this technique can in principle be applied to analyse arbitrarily-shaped regions on the celestial sphere. The assumed isotropy of the foreground spectral parameters over the sky is another extreme simplification. For example, the synchrotron spectral index varies in a wide range. Giardino \\et (2002) calculated the synchrotron temperature spectral index using three low-frequency radio survey maps and found it to vary in the range 2.5$\\le \\beta_{408/1420} \\le$3.2. Variation of the dust colour temperature $T_{\\rm dust}$ is more important for the {\\sc Planck} experiment because the HFI channels are quite sensitive to thermal dust emission. Schlegel \\et (1998) found $T_{\\rm dust}$ to vary in the range 16K to 20K, about a mean value $\\langle T_{\\rm dust} \\rangle$=18K assuming a single--component model. The effect of a spatially-varying Galactic dust emissivity index $\\beta$ on the MEM reconstruction was investigated for case of a flat-sky patch by Jones \\et (2000). Several dust sub-components with different emissivities, but with the same colour temperature, were included in the separation process, which provided a good reconstruction of the components. More recently, Barreiro \\et (2004) used a combined real and harmonic space-based MEM technique to perform a component separation on real data, in the presence of anisotropic noise, cut-sky maps and spectral index uncertainties. Since this approach requires multiple transitions between pixel and spectral domains, however, the computation of the necessary spherical harmonic transforms makes it is much slower than harmonic-space MEM, and hence it was implemented only for low-resolution COBE data. In this paper we will demonstrate how to extend the full-sky harmonic-space MEM component separation method to take into account anisotropic noise and variations in spectral parameters, by making use of prior knowledge of the uneven sky coverage and the average value of the spectral parameters. The structure of the paper is as follows. In a Section~\\ref{noise} we summarise the basics of the MEM component separation technique, describe the model of the microwave sky used in the simulations, and review the impact of the non-isotropic noise on the CMB reconstruction. In a Section~\\ref{dust_temp} we will introduce an approach for taking account of dust colour temperature variations in the component separation process, describe the new microwave sky model with variable $T_{\\rm dust}$ and show the results for the reconstruction tests made using different approximations. In Section~\\ref{discussion} we discuss the results and present our conclusions. ", "conclusions": "\\label{discussion} In this paper, we have demonstrated an approach that allows one to account for the spatial variations of the noise properties and spectral characteristics of foregrounds in the harmonic-space maximum-entropy component separation technique. In Section~\\ref{results_rms}, we show that the impact of a realistic level of the anisotropic noise on the quality of the foreground separation is quite small, at least for the simple scanning strategy assumed in generating our simulated observations. This can be explained by the fact that the rms noise level differs from the average level only in relatively small areas around ecliptic poles. This leads to variations in the noise level at each harmonic mode that differ from the average value by only a few per cent. We also illustrate in Section~\\ref{results_rms} that the it is possible to perform harmonic-space component separation on cut-sky maps by treating the cut as an extreme example of anisotropic noise. This appraoch has the advantage of not requiring one to smooth the edges of the cut with some apodising function prior to the analysis. In Section~\\ref{dust_temp}, we show that the variation of spectral parameters of foregrounds may be taken into account by a method of succesive approximations based on a series expansion of the corresponding intensity field around the mean value of the parameter. In particular, we investigate the effect of dust colour temperature variations on the quality of the component separation, focussing on the reconstruction of the CMB. We show that realistic dust temperature variations lead to severe contaminaton of the CMB reconstruction if, in the separation process, the dust temperature is assumed not to vary. This contamination is concentrated in the Galactic plane, but significant artefacts exist at high Galactic latitudes. The poor quality of the reconstruction outside the Galactic plane is a result of performing the reconstruction mode-by mode in harmonic space. The inaccurate model of the dust emission leads to errors in the determined amplitudes of a wide range of spherical harmonics in the CMB reconstruction. Many of these modes do not lie predominatly in the Galactic plane region, but contribute to the reconstruction over the whole sky. If one is content simply with removing foregrounds from the CMB, rather than performing a component separation, then one could apply a Galactic cut prior to the analysis. In Fig.~\\ref{cmb_dust_cut}, we show the CMB reconstruction and residuals respectively obtained by applying a Galactic cut of $\\pm 25\\degr$ to each simulated frequency map used in Section~\\ref{model_temp}, and assuming a constant dust temperature across the sky. \\begin{figure*} \\begin{center} \\centerline{\\epsfig{file=CMB_map_restored_1024_mkK_cut25.ps,width=8.5cm} \\quad \\epsfig{file=CMB_map_resid_1024_mkK_cut25.ps,width=8.5cm}} \\caption{The reconstruction (left) and residuals (right) of the CMB for the case plotted in Fig.~\\ref{cmb_recon} (upper row), but with a Galactic cut of $\\pm 25\\degr$. The maps are plotted in units of $\\mu K$.} \\label{cmb_dust_cut} \\end{center} \\end{figure*} The noise rms in the cut region was assumed to be formally infinite in the manner discussed in Section~\\ref{noise}. We see from the figures that the quality of the reconstruction is significantly improved as compared with the case in which no Galactic cut was applied, which was shown in Fig.~\\ref{cmb_recon} (top row). In particular, we note that the residuals now contain no obvious artefacts outside of the cut region. Thus, even assuming an inaccurate dust model, one can still recover a reasonable reconstruction of the CMB outside of the Galactic plane. It is clearly not possible, however, to perform an acceptable all-sky component separation for the {\\sc Planck} experiment by assuming constant dust spectral parameters, even using all 9 frequency channels. Nevertheless, taking account of dust temperature variations up to first order in the series expansion significantly improves the CMB reconstruction to an acceptable level. This reconstruction quality is still further improved by including second-order terms and is then comparable to that obtained for the ideal case presented in Stolyarov \\et (2002), where the simulated observations assumed no dust temperature variation across the sky. Moreover, an accurate reconstruction of the dust temperature variation is obtained over the whole sky. Finally we note that the approach for dealing with spatially-varying spectral parameters described in this paper can also be applied to other foregrounds to yield, for example, maps of the synchrotron spectral index. The method can also be used to reconstruct the electron temperature in clusters from their thermal SZ effect, as will be discussed in a forthcoming paper." }, "0405/astro-ph0405177_arXiv.txt": { "abstract": "We report on submillimetre (submm) observations of three high redshift quasars ($z>$6) made using the SCUBA camera on the James Clerk Maxwell Telescope (JCMT). Only one of the sample was detected ($>10\\sigma$ significance) at 850$\\mu m$ -- SDSS J1148$+$5251 ($z=6.43$). It was also detected at 450$\\mu m$ ($>3\\sigma$ significance), one of the few quasars at $z>4$ for which this has been the case. In combination with existing millimetric data, the 850$\\mu$m and 450$\\mu$m detections allow us to place limits on the temperature of the submm-emitting dust. The dust temperature is of no trivial importance given the high redshift of the source, since a cold temperature would signify a large mass of dust to be synthesized in the little time available (as an extreme upper limit in only 0.9Gyr since $z=\\infty$). We find, however, that the combined millimetre and submm data for the source cannot simply be characterised using the single-temperature greybody fit that has been used at lower redshifts. We discuss the results of the observing and modelling, and speculate as to the origin of the deviations. ", "introduction": "A number of independent lines of investigation over the last 10 years have placed submillimetre (submm) observations of high-redshift quasars into the spotlight. Observations have unveiled a population of extremely luminous submm sources lying at high redshift, believed to be the dust-obscured, star-forming ancestors of massive elliptical galaxies. Contemporaneously, it was realised that a tight correlation exists between the stellar velocity dispersion of galactic spheroids, and the mass of their central, supermassive black hole (Gebhardt et al. 2000). Taken together, these indicate that luminous active galactic nuclei (AGN) at high redshift--- the build-up phase of a supermassive black hole --- are prime sites at which to search for the dust-enshrouded star-burst phase through which, according to the new galaxy-formation paradigm, their massive spheroids necessarily must pass. The high, sustained luminosity of quasars across the electro-magnetic spectrum, allows them to be studied over a wide range of both redshift and observing wavelength. Early observations of z$\\sim$4.5 quasars by McMahon et al. (1994), and Isaak et al. (1994) established that some high redshift quasars were prodigious far infrared emitters with $\\rm L_{fir}\\sim10^{13-14} L_{\\odot}$ and estimated masses of cool dust of $\\rm\\sim 10^{8-9} M_{\\odot}$. The discovery of quasars at $z>6$ (Fan et al. 2003) now makes it possible to compile homogeneous, well-defined samples over a significant span of the lifetime of the cosmos, from recent times to the threshold of reionization. Follow-up is simplified by the accurately-known optical positions and the spectroscopic redshifts of the host galaxies, which can readily be determined to the precision required to pinpoint emission lines from molecular gas--- a key indicator of the conditions required for star formation. Recent SCUBA studies of the submm emission from high redshift ($z>4$), radio-quiet quasars have been reported by McMahon et al. (1999), Isaak et al. (2002) and Priddey et al. (2003b), along with a sample at lower redshift (z$\\sim$2) by Priddey et al. (2003a). A considerable fraction of the targets have been shown to be luminous submm sources. Interestingly, this fraction appears to have no significant dependence upon redshift. Similar conclusions have been drawn from observations at millimetre (mm) wavelengths (eg. Omont et al. (2001), Carilli et al. (2001)). ", "conclusions": "The rest-frame FIR spectral energy distribution is key to determining the thermal origin of the observed submm emission from high-redshift quasars. The observations presented here suggest that the host galaxies of quasars out to redshifts of $z>6$ are actively undergoing star formation. Multi-wavelength observations spanning the mm and submm are crucial to providing constraints on the dust mass, far-infrared luminosity and inferred star formation rate, thus further exploring the role of star formation in high-redshift quasar host galaxies." }, "0405/gr-qc0405130_arXiv.txt": { "abstract": "\\ \\\\ We identify conditions for the presence of negative specific heat in non-relativistic self-gravitating systems and similar systems of attracting particles. The method used, is to analyse the Virial theorem and two soluble models of systems of attracting particles, and to map the sign of the specific heat for different combinations of the number of spatial dimensions of the system, $D$($\\geq 2$), and the exponent, $\\nu$($\\neq 0$), in the force potential, $\\phi=Cr^\\nu$. Negative specific heat in such systems is found to be present exactly for $\\nu=-1$, at least for $D \\geq 3$. For many combinations of $D$ and $\\nu$ representing long-range forces, the specific heat is positive or zero, for both models and the Virial theorem. Hence negative specific heat is not caused by long-range forces as such. We also find that negative specific heat appears when $\\nu$ is negative, and there is no singular point in a certain density distribution. A possible mechanism behind this is suggested. ", "introduction": "In the area of self-gravitating systems, the literature often gives the impression that the long-range properties of forces are the cause of negative specific heat in such systems. Here, we will show that only some systems affected by long-range forces exhibit negative specific heat. We limit our investigations to potentials of the form \\begin{equation} \\label{188} \\phi(r) = C r^\\nu \\end{equation} and of the form \\begin{equation} \\label{301} \\phi(r) = C_1 \\ln{C_2 r} \\end{equation} where $r$ is the radial coordinate in a space with $D$ dimensions ($D \\geq 2$), $\\nu$ an integer ($\\nu \\neq 0$), and $C$, $C_1$ and $C_2$ constants (only depending on $D$ and $\\nu$). Observe that gravitation is described by (\\ref{188}) for $D \\neq 2$, and by (\\ref{301}) for $D=2$. We will, however, to extend the analysis, regard both (\\ref{188}) and (\\ref{301}) for any $D \\geq 2$. We can use spaces with $D \\geq 2$ and $\\nu \\neq 0$ as our theoretical \"test bench\" to investigate the cause of negative specific heat. There is, for instance, no a priori reason why long-range forces should result in negative specific heat in three dimensions, but not in other number of dimensions. If this happens only in three dimensions, there is probably another mechanism laying behind, that is more relevant to explain negative specific heat. Before investigating negative specific heat, we will define the concept of long- range and short-range forces. Let us regard a continuous medium with constant density, and investigate from which areas of the medium the main part of the potential energy of a particle embedded in the medium comes. Short-range forces are, for potentials of the form (\\ref{188}), then characterised by \\begin{equation} \\label{190} D + \\nu < 0 \\end{equation} and long-range forces by \\begin{equation} \\label{191} D + \\nu > 0 \\end{equation} In the case where $D+\\nu=0$, the dependence of distance for the total potential energy of the particle is logarithmic, and the force is then both short- and long-range. For potentials of the form (\\ref{301}), the corresponding criteria implies that we have a force that is both short- and long-range. The definition of specific heat at constant volume is \\begin{equation} \\label{192} C_V \\equiv \\frac{\\partial }{\\partial T} \\vert_V \\end{equation} where \"$<>$\" expresses a time average. Negative specific heat was first investigated by Lynden-Bell and Wood~\\cite{Lynd3} and Thirring~\\cite{Thir}. It is also described by Hut~\\cite{Hut}, Lynden-Bell~\\cite{Lynd1}, Lynden-Bell and Lynden-Bell~\\cite{Lynd2} and Padmanabhan~\\cite{Padm}. The Virial theorem in its simplest form, without external pressure, applied to a system of particles interacting with potentials of the form (\\ref{188}) reads \\begin{equation} \\label{193} = \\frac{\\nu + 2}{\\nu} \\end{equation} Using (\\ref{192}) and the assumption $\\frac{\\partial }{\\partial T} \\vert_V > 0$, which is valid for many systems, we obtain that negative specific heat appears exactly for $\\nu=-1$, independent of the number of dimensions, $D$($\\geq 1$), of the system. In three dimensions $\\nu \\geq -2$ corresponds to long-range forces. So, for three-dimensional systems (and negative $\\nu$), negative specific heat seems to be correlated with the long-range nature of the forces. The Padmanabhan model and the Lynden-Bell model (see below) were invented to verify that $\\nu=-1$ gives negative specific heat in three-dimensional systems. When considering systems with $D>3$, it becomes apparent that the domain (in $D$ and $\\nu$) of negative specific heat predicted by the Virial theorem, is just a small part of the domain with long-range forces. We will show that generalised versions of the Padmanabhan model and the Lynden-Bell model (both with $D \\geq 2$) give the same domain of negative specific heat as the Virial theorem, for $D \\geq 3$ and any $\\nu$($\\neq 0$). For systems with particles interacting with potentials of the form (\\ref{301}), and without external pressure, the Virial theorem reads \\begin{equation} \\label{302} = \\mathrm{constant} \\end{equation} That is, energy supplied only contributes to the potential energy of the system. Here we clearly have $\\frac{\\partial }{\\partial T} \\vert_V = 0$, so this is one of those systems where the spatial volume is important in the contribution to the temperature. We will show that, for this kind of system, the two generalized models give $C_V > 0$ for all $D$($\\geq 2$). The result of the analyses of systems with potentials of the form (\\ref{188}) is mapped in fig.~1. The border between long- and short-range forces is marked. Gravitation obeys, for $D \\neq 2$, $\\nu = 2 - D$, and is therefore, according to (\\ref{191}), a long-range force for $D \\neq 2$. Only for $\\nu = -1$ we have significant negative specific heat in our models. Only the Padmanabhan model has a non-negative value on $C_V$ for $\\nu = -1$, namely when $D = 2$. It is, however, as often suggested, reasonable to believe that long-range forces in a system is the cause of non-extensivity of the system. With non-extensivity, we mean that if two originally separated subsystems with energy $E_1$ and $E_2$ are combined, their total energy will not be $E_1 + E_2$. For systems defined by short-range forces, the interaction energy between subsystems becomes negligible for large enough subsystems, since the interaction energy scales as the area of the subsystems, and then we have an extensive system. For subsystems interacting by long-range forces, the interaction energy remains significant, since the interaction energy scales as the volume of the subsystems, and we then have a non-extensive system. Oppenheim~\\cite{Oppe} investigates the statistical physics of systems with long-range interactions. ", "conclusions": "For systems consisting of particles interacting with potentials of the form (\\ref{188}), the Padmanabhan model and the Lynden-Bell model give the same interval (in $D$ and $\\nu$) of negative specific heat as the Virial theorem does, for $D \\geq 3$. Significant negative specific heat is present for $\\nu=-1$. Long-range forces obeys $D+\\nu>0$. Hence, the long-range nature of a force is not sufficient to give negative specific heat. To use the Virial theorem to identify negative specific heat in systems with arbitrary interaction potentials, the radial density (or probability) distribution has to be regarded. To derive these distributions, and their relation to energy and temperature, is crucial to identify negative specific heat in these systems, if the Virial theorem will be used. A physical mechanism behind negative specific heat, independent of the Virial theorem, is suggested. It implies that positive specific heat appears when the system is strongly affected by what we call walls of the first kind, or the density distribution shows that many particles are affected by what we call walls of the second kind. In all other cases negative specific heat is assumed to be present. This is consistent with the results of the two soluble models, except for in one case (the Padmanabhan model with $D=2$ and $\\nu=-1$). The principle may, in some developed form, be an important part of \"the general cause to negative specific heat\"." }, "0405/astro-ph0405388_arXiv.txt": { "abstract": "{We present results of first simultaneous optical and X-ray observations of peculiar binary system SS433. For the first time, chaotic variability of SS433 in the optical spectral band ($R$ band) on time scales as small as tens of seconds was detected. We find that the X-ray flux of SS433 is delayed with respect to the optical emission by approximately 80 sec. Such a delay can be interpreted as the travel time of mass accretion rate perturbations from the jet base to the observed X-ray emitting region. In this model, the length of the supercritical accretion disk funnel in SS433 is $\\sim 10^{12}$ cm. ", "introduction": "SS433 is the unique X-ray binary, demonstrating precessing jets and supercritical accretion disk around a black hole (Margon, 1984; see Fabrika, 2004 for a recent review). The total luminosity of the system is estimated to be at a level of $10^{40}$ erg/s, with most of this luminosity coming in UV and optical bands (Cherepashchuk et al., 1982; Dolan et al., 1997). According to the current model of this binary, the dense outflowing wind and the geometrically thick accretion disk totally cover from us the direct X-ray emission from the innermost hot regions of the accretion disk. The early observations of SS433 by EXOSAT and {\\it ASCA} satellites led to the conclusion that standard X-ray emission in SS433 originates in optically thin plasma outflowing in a mildly relativistic, $v=0.26c$, jets (Watson et al. 1986, Kotani et al. 1996). A detailed analysis of high-resolution {\\it Chandra} observations of SS433 suggests that the X-ray jet consists of well collimated freely expanding cooling plasma with temperature varying from 20 to 0.5 keV. The size of the vizible X-ray jet is estimated to be $\\sim 10^{10}-10^{11}$ cm (Marshall et al. 2002). This is the lower limit on the size of X-ray emitting jet if heating effects are not important (Brinkmann and Kawai 2000). The observed X-ray luminosity ($\\approx 10^{36}$ erg/s) is only a tiny part of the total jet power. For example, the kinetic power of the jet is estimated to be 1000 times larger than the observed X-ray luminosity. The contribution of the observed X-ray luminosity to the total energetics of the system is even smaller -- not more than $10^{-4}$. The long-term behavior of SS433 has been well studied at different wavelengths. The source demonstrates different types of long-term periodicities: precessional (162.3 days), orbital (13.082 days), and nutational (6.28 days) ones. SS433 exhibits noticeable variability on shorter times as well, however at time scales smaller than hours this has been poorly studied. For example, erratic variability with an amplitude of 5-10\\,\\% on a time-scale of a few minutes was found in the optical band by Goranskii et al. (1987) and Zwitter et al. (1991). Recently, Kotani et al. (2002) detected X-ray variability of SS433 (RXTE, PCA) on a time scale of $\\sim 50$~s, when the source was in its active state. The accretion disk can modulate energy release and luminosity in the broad range of time scales (\\cite{lyubarsky97}). This mechanism can nicely explain the observed flicker-noise spectra found in some X-ray binaries, e.g. Cyg X-1 (\\cite{chur01}). In the case of SS433 the X-ray emission of the accretion disk is not directly visible, so variability in the optical and UV may be a better indicator of the temporal properties of the accretion disk. The UV and optical radiation of SS433 is well approximated by a single reddened black body source (Cherepashchuk et al., 1982; Dolan et al., 1997; Gies et al., 2002) with a temperature $T_e \\sim 5 \\cdot 10^4$~K ($A_V \\approx 8$) and a size $r_{ph} \\sim (1-2) \\cdot 10^{12}$~cm. The bulk of the optical and UV emission most likely escapes from the hot funnel in the photosphere close to the jet bases (Dolan et al. 1997; Fabrika 2004). Should X-ray and optical variabilities be due to one physical reason (e.g., the accretion rate modulation), simultaneous observations of SS433 in the optical band (which reflect variability of the energy release in the supercritical accretion disk) and in X-rays (that come from the footpoint of the jet) will be invaluable to test the origin of jets in SS433 and their close surroundings. In this paper we present the analysis of simultaneous optical and X-ray observations of SS433 in March 2004. ", "conclusions": "Simultaneous observations of SS433 in optical and X-ray energy bands demonstrate clear correlation of measured fluxes. The variable part of the X-ray lightcurve is delayed with respect to the optical one by about 80 sec. The lag of X-rays with respect to the optical variability can be anticipated in the framework of the cooling jet model that appears as a result of acceleration of matter in the base of the funnel in the center of the supercritically accrettion disk (Calvani, Nobili 1981; Bodo et al. 1985; Eggum et al. 1988; Okuda 2002). The variability of SS433 on longer time scales with lagged correlation has been known in different parts of the emitting regions of this system (Cherepaschuk 2002, Fabrika 2004). For example, the optical nutation variability has been observed to preceed the nutation motion of optical jets by 0.6 days. This is exactly the time it takes for the jet gas to travel from the compact object to the $H\\alpha$ emitting region downstream the jets. A maximum of $H\\alpha$ emission in the optical jets forms at a distance of $4 \\cdot 10^{14}$cm from the source (Borisov and Fabrika 1987). The amplitude of the nutation variability (Cherepaschuk 2002) is $\\Delta V \\approx 0.17$ and increases in shorter $B$ band. This nutational lag proves that the optical and UV radiation come from the inner wind region, i.e. from the jet base, but not from the jet itself. The jets' energy budget ($L_{kin}\\sim 10^{39}$~erg/s, $L_{x}\\sim 10^{36}$~erg/s) is insufficient to provide the optical nutational modulation, which is 20~\\% of $L_{opt}\\sim 10^{38}$~erg/s, $L_{UV}\\sim 10^{40}$~erg/s. In our case we are dealing with two nearby emitting regions - the innermost region of the supercritical accretion disk, which emits most of the energy in the optical and UV band, and the X-ray emitting jet, which is supposed to be launched near the compact object and have a bulk velocity of $0.26c$. The most straightforward and simplest interpretation of the observed time delay between the optical and X-ray fluxes, which does not contradict the common knowledge about SS433, can be as follows. Mass accretion fluctuations appear in the bottom of the supercritical disk funnel where the jet is being formed and results in fluctiating X-ray flux inside the funnel. According to the supercritical disk simulations (Eggum et al. 1988; Okuda 2002), the place of the jet formation (10-100~$r_g$) is quite close to the black hole. The X-ray radiation coming upwards the funnel may be absorbed by the outer funnel walls (which are directly observed or lie immediately close to the directly observed parts of the funnel) and produce fluctuating UV-optical radiation via photoabsorption. Similar processes are well-known to operate in low mass X-ray binaries (see e.g. \\cite{lawrence83}). The same fluctuations reach the observed part of the X-ray jet after the jet propagation time inside the funnel and may appear as variable fraction of the visible X-ray emission. This is the basic mechanism we propose to explain the correlated X-ray - optical variability. It is important that this correlated variability between optical and X-ray variations has to be related to all time-scales generated in the accretion disk, because we do not see a QPO-like variability. The delayed X-ray variability allows us to estimate the funnel length $r_f$ in SS433. The time it takes for the X-ray jet to appear above the funnel edge is $r_f/V_j$, the time for the light to reproduce optical emission via photoabsorption of X-rays from the funnel bottom is $\\eta r_f/c$, where factor $\\eta \\ga 1$ is responsible for possible delays (geometry, scattering, etc.). The difference between these two times yields the time delay between the optical and X-ray variabilities found by us, $\\Delta t \\sim 80$~sec. So for the funnel length we have the estimate $r_f \\sim 2 \\cdot 10^{12}/(c/V_j-\\eta)$~cm. A lower limit on the funnel length is for $\\eta=1$, i.e. $r_f \\ga 10^{12}$~cm. So the bases of the X-ray jets, whose length $\\ga 10^{10}-10^{11}$~cm were derived from the {\\it Chandra} spectra (Marshall et al. 2002), have to be placed at distance $r_f\\sim 10^{12}$ cm from the black hole." }, "0405/astro-ph0405341_arXiv.txt": { "abstract": "An extremely cold and big spot in the WMAP 1-year data is analyzed. Our work is a continuation of a previous paper (Vielva et al. 2004) where non-Gaussianity was detected, with a method based on the Spherical Mexican Hat Wavelet (SMHW) technique. We study the spots at different thresholds on the SMHW coefficient maps, considering six estimators, namely number of maxima, number of minima, number of hot and cold spots, and number of pixels of the spots. At SMHW scales around $4^\\circ$ ($10^\\circ$ on the sky), the data deviate from Gaussianity. The analysis is performed on all sky, the northern and southern hemispheres, and on four regions covering all the sky. A cold spot at ($b = -57^\\circ, l = 209^\\circ$) is found to be the source of this non-Gaussian signature. We compare the spots of our data with 10000 Gaussian simulations, and conclude that only around 0.2\\% of them present such a cold spot. Excluding this spot, the remaining map is compatible with Gaussianity and even the excess of kurtosis in Vielva et al. 2004, is found to be due exclusively to this spot. Finally, we study whether the spot causing the observed deviation from Gaussianity could be generated by systematics or foregrounds. None of them seem to be responsible for the non-Gaussian detection. ", "introduction": "The Cosmic Microwave Background (CMB) is the most ancient image of the universe. The measured temperature fluctuations reveal important characteristics of the primitive universe. Nowadays many models try to explain the primordial density fluctuations, and the key for deciding among them is the Gaussianity of the temperature fluctuations of the CMB. According to standard inflationary theories, these fluctuations are generated by a Gaussian, homogeneous and isotropic random field, whereas non-standard inflation (Linde \\& Mukhanov 1997, Peebles 1997, Bernardeau \\& Uzan 2002 and Acquaviva et al. 2003) and topological defect models (Turok \\& Spergel 1990 and Durrer 1999), predict non-Gaussian random fields. Many difficulties have to be avoided in the Gaussianity analysis, since the temperature fluctuations may come from different sources, which should be identified. The primary fluctuations were originated on the last scattering surface (LSS), while secondary fluctuations were produced in the trip that CMB photons cover from the LSS to us. Some examples of the latter are the fluctuations caused by the reionization of the universe (Ostriker \\& Vishniac 1986 and Aghanim et al. 1996), the Rees-Sciama effect (Rees \\& Sciama 1968 and Mart{\\'\\i}nez--Gonz{\\'a}lez \\& Sanz 1990), or gravitational lensing (Mart{\\'\\i}nez--Gonz{\\'a}lez et al. 1997, Hu 2000 and Goldberg \\& Spergel 1999). In addition we have systematic effects and contaminating foregrounds, like the Galactic emissions or extragalactic point sources, which have to be considered in the Gaussianity study. In the last years, many groups have published analysis of Gaussianity. The Wilkinson Microwave Anisotropy Probe (WMAP) 1-year all-sky data (Bennett et al. 2003a), were found to be consistent with Gaussianity by the WMAP science team (Komatsu et al. 2003). A measure of the phase correlations of temperature fluctuations (taking several combinations of the bispectrum into account), and the Minkowski functionals, were used in the latter paper. Thereafter several other groups have found non-Gaussian signatures or asymmetries in the WMAP data (Park 2003, Eriksen et al. 2003, 2004a, Hansen et al. 2004a, b). Several methods were used including Minkowski functionals, \\emph{N}-point correlation functions, or local curvature methods. Also some other works such as Chiang et al. 2003 and Copi et al. 2003, detect deviations from Gaussianity, using the Internal Linear Combination (ILC) map (Bennett et al. 2003a) and other maps, derived from the WMAP data. (However, the use of the ILC map is not recommended by the WMAP team, because of their complex noise and foreground properties). Our basic reference is Vielva et al. 2004, where a non-Gaussian detection in the WMAP 1-year data was reported. Convolving the data with the SMHW, an excess of kurtosis was found at scales around $4^\\circ$, involving a size on the sky close to $10^\\circ$. The deviation from Gaussianity presented an upper tail probability of 0.4\\%. This deviation was localized in the southern hemisphere and an extremely cold spot at $b = -57^\\circ, l = 209^\\circ$, was regarded as a possible source of the non-Gaussianity (we call it \\emph{The Spot}). Mukherjee \\& Wang 2004 have carried out an extrema analysis using the SMHW, finding an excess of cold pixels in all the sky at the same scales as Vielva et al. 2004. In the present paper, our aim is to localize the non-Gaussianity, specifying whether The Spot alone is the origin of the detection, and to study all the spots of the data in order to quantify how atypical The Spot under a Gaussian model is. Also recently, Larson \\& Wandelt 2004 have carried out an extrema analysis in real space, finding non-Gaussian evidences. The present paper is organized as follows. In $\\S$2 we describe the data and simulations used to study the Gaussianity. The wavelet technique and the estimators are explained in $\\S$3. The results of our work are presented in $\\S$4. In $\\S$5 we discuss the possible sources of the non-Gaussian detection and in $\\S$6 we present the conclusions of our results. ", "conclusions": "Motivated by the non-Gaussianity found in the WMAP 1-year data using the SMHW, we have performed an analysis of the spots in the SMHW coefficients map, aimed to locate possible contributors in the sky. An extremely cold and big spot is detected. This spot (\\emph{The Spot}) is seen in the SMHW coefficients at scales around $4^\\circ$ (implying a size of around $10^\\circ$ on the sky) and at Galactic coordinates $b=-57^\\circ, l=209^\\circ$. The probability of having such spot for a Gaussian model is of only $\\approx 0.2\\%$, which implies that, if intrinsic, The Spot has not been originated by primary anisotropies in the standard scenario of structure formation since standard inflation predicts Gaussian fluctuations in the matter energy density and therefore in the CMB temperature fluctuations. When this spot is not considered in the analysis the rest of the data seem to be consistent with Gaussianity. In order to identify the source of The Spot we have performed several tests related to systematic effects and foregrounds. We have checked that uncertainties in the noise or in the beam response have a negligible effect in our results at the relevant wavelet scales. Looking at the maps corresponding to the different receivers, we see a clear consistency in the area, amplitude and position of The Spot. Hence our detection is not due to any deficient receiver. In relation to the possible foregrounds contribution, we have looked for possible frequency dependences in the amplitude and area of The Spot. Again both quantities show a nice consistency with a constant line in the range from 23 to 94 GHz. Whereas the Galactic foregrounds show a very different frequency dependence with respect to the constant behaviour, the SZ effect does not separate much from it in that frequency range. A comparable spot could be produced either by the Coma cluster at a much closer distance, or by several rich clusters at the actual distance of Coma. We have checked that no nearby rich cluster of galaxies is located in the position of The Spot. Finally, intrinsic fluctuations cannot be rejected as the source of The Spot. In particular, a massive and distant super-structure could in principle produce a decrement as the one observed through the Rees-Sciama effect (Mart\\'\\i nez-Gonz\\'alez and Sanz 1990). This massive structure (of order of at least $10^{16} M_\\odot$) should be placed far enough because otherwise it would have been detected previously. Alternatively, more speculative possibilities are topological defects (monopoles or textures) or non-standard inflationary scenarios. Even more, a combination of secondary anisotropies with primary ones, cannot be rejected as the source of our non-Gaussian spot. For instance a possibility could be a combination of the Sunyaev-Zeldovich effect with a Sachs-Wolfe plateau." }, "0405/astro-ph0405355_arXiv.txt": { "abstract": "Motivated by the WMAP results indicating an early epoch of reionization, we consider alternative cosmic star formation models which are capable of reionizing the early intergalactic medium. We develop models which include an early burst of massive stars (with several possible mass ranges) combined with standard star formation. We compute the stellar ionizing flux of photons and we track the nucleosynthetic yields for several elements\\,: D, $^{4}$He, C, N, O, Si, S, Fe, Zn. We compute the subsequent chemical evolution as a function of redshift, both in the intergalactic medium and in the interstellar medium of forming galaxies, starting with the primordial objects which are responsible for the reionization. We apply constraints from the observed abundances in the Lyman $\\alpha$ forest and in Damped Lyman $\\alpha$ clouds in conjunction with the ability of the models to produce the required degree of reionization. We also consider possible constraints associated with the observations of the two extremely metal-poor stars HE 0107-5240 and CS22949-037. We confirm that an early top-heavy stellar component is required, as a standard star formation model is unable to reionize the early Universe and reproduce the abundances of the very metal-poor halo stars. A bimodal (or top-heavy) IMF (40 - 100 M$_\\odot$) is our preferred scenario compared to the extreme mass range ($\\ga$ 100 M$_\\odot$) often assumed to be responsible for the early stages of reionization. A mode of even more extreme stellar masses in the range ($\\ge$ 270 M$_{\\odot}$) has also been considered. All massive stars in this mode collapse entirely into black holes, and as a consequence, chemical evolution and reionization are de-correlated. The ionizing flux from these very massive stars can easily reionize the Universe at $z\\sim 17$. However the chemical evolution in this case is exactly the same as in the standard star formation model, and the observed high redshift abundances are not reproduced. We show that the initial top-heavy mode, which originally was introduced to reionize the early Universe, produces rapid initial metal pollution. The existence of old, C-rich halo stars with high [O/Fe] and [C/Fe] ratios is predicted as a consequence of these massive stars. The recently observed abundances in the oldest halo stars could trace this very specific stellar population. The extreme mass range is disfavored and there is no evidence, nor any need, for a hypothesised primordial population of very massive % stars in order to account for the chemical abundances of extremely metal-poor halo stars or of the intergalactic medium. The combined population of early-forming, normal (0.1 - 100 M$_\\odot$) and massive (40 - 100 M$_\\odot$) stars can simultaneously explain the cosmic chemical evolution and the observations of extremely metal-poor halo stars and also account for early cosmological reionization. ", "introduction": "With the recent release of WMAP data, we have without question, entered the age of precision cosmology. While many of the accurately determined cosmological parameters were found to have values consistent with a $\\Lambda$CDM cosmology, the $TE$ cross-correlation power spectrum is consistent with a large optical depth due to electron scattering, $\\tau_e = 0.17 \\pm 0.04$ \\citep{kogut:03} which suggests a very early epoch ($z \\simeq 20$) of reionization in the inter-galactic medium (IGM). It is often suggested that reionization is due to a population of very massive stars forming in minihalos, cooled predominantly by H$_2$ \\citep[see][]{cen:03a,haiman:03,wyithe:03,bromm:04}. It has also been suggested that this early period of reionization is not total, but rather is regulated by radiative feedback effects delaying the complete reionization to a later redshift ($z \\simeq 6$) accounting for the observed Gunn-Peterson trough \\citep{becker:01}. Even a relatively brief period of massive star formation at high redshift can have dramatic consequences on the chemical enrichment of the primitive structures and of the IGM. There has been significant progress in our understanding of the yields of massive stars, particularly for so-called pair-instability supernovae, allowing one to trace the chemical history as influenced by an early massive population of stars. If indeed the first generation of stars were very massive (140-260 M$_\\odot$), they would act as prompt initial enrichment sources for the IGM as well as the progenitors of galaxies such as our own. In addition, they would lay down a very distinct chemical signature which can be compared with other potential initial mass functions for the first generation of stars capable of producing the required degree of reionization. The argument for a prompt initial enrichment laying down the seeds for the chemical evolution of the Galaxy is an old one going back to \\citet*{truran:71}. This notion receives confirmation in observations of r-process elements as a function of [Fe/H] in halo stars \\citep*{wasserburg:00,qian:01}. The question of the ionization capacity of the massive stars supposed to produce this initial enrichment was recently considered by \\citet{oh:01}, \\citet{oh:02}, and \\citet*{venkatesan:03b} in the context of a late period of reionization at a redshift $z \\sim 6$. More recently, metallicity studies of Damped Lyman $\\alpha$ clouds (DLAs) at high redshift also indicate the presence of an initial metallicity of roughly [Fe/H] $\\sim 10^{-3}$ \\citep{prochaska:03}. \\citet*{tumlinson:04} examined this question in the context of an early period of reionization as indicated by the WMAP data and present a general discussion of the initial mass function (IMF) of the first stars and their associated nucleosynthesis. Specifically they use the measured element abundances in the metal-poor halo stars to derive the IMF of population III stars responsible for the reionization of the early IGM. While our goal here is similar, we include a comprehensive model of chemical evolution to address this question. Related to the question of an early population (Pop III) generation of stars is the observation of extremely metal-poor ([Fe/H] $<$ -4) stars such as HE 0107-5240 (\\citealt{christlieb:04} and \\citealt*{bessell:04}), a halo giant which is extremely metal-poor, [Fe/H] = -5.3. While the abundances of other observed elements, e.g., Na, Mg, Ca, etc. are (roughly) as low, the abundance of C (and N and O) is surprisingly high, [C/H] $\\approx -1.3$ ([N/H] $\\approx -3.0$, [O/H] $\\approx -2.9$). Indeed the very existence of this star poses a challenge for both models of chemical evolution (and also stellar evolution) and star formation as its mass is believed to be quite low ($M \\sim 0.8$ M$_\\odot$). We note, however, that it is possible that this star (and other similar stars such as CS 22949-037, \\citealt{depagne:02} and \\citealt{israelian:04}) was born under very particular circumstances as these abundance patterns may be explained by the pre-enrichment of a massive zero-metallicity star (with fall-back) (\\citealt*{umeda:03}; see also \\citealt{suda:04}), or perhaps the star was polluted by a binary companion in its AGB phase. Thus with some degree of caution, we will discuss the implications of the existence of this star with respect to the models considered. While there have been several detailed studies concerning the ionization efficiency of massive stars, attention to chemical enrichment has focused on the overall metallicity produced in structures with massive or very massive stars. Here, we calculate the combined effect of an early population of massive stars on the reionization history and chemical enrichment history of the early cosmic structures and of the IGM. In the framework of the hierarchical galaxy formation model, most of the halo, and in particular a DLA, is built up by smaller merging systems that are tidally disrupted (see e.g. \\citep{zentner:03}). Therefore, they inherit the metals produced at large z by population III stars as a Prompt Initial Enrichment (PIE). The intensity of this PIE must be limited to avoid a metallicity surpluss in these late structures which produce new metals via the normal mode of star formation. Thus our approach is the following : (i) since the metallicity increases very rapidly as soon as the first population III stars explode, the timescale of the massive mode of star formation is necessarily short (see section 4.3); (ii) therefore, the cosmic SFR at $z \\la 5-6$ reflects only the normal mode of star formation. This allows one to normalize its intensity to fit the observed cosmic SFR at these redshifts ; (iii) once this normalization has been done, the metal production by the normal mode is immediately derived without any other parameter adjustement; (iv) from the amount of metals produced by the normal mode, it is then possible to estimate the maximum permitted initial metallicity in the structures at the beginning of the normal mode to avoid the over-production of metals in the late structures; (v) finally, this allows us to normalize the intensity of the early massive mode (pop. III stars) as it governs the PIE. We consider several initial mass functions which differ primarily in the population of massive stars forming at zero (or near zero) metallicity. Given an initial mass function and an associated star formation rate, we calculate in a consistent way both the chemical history of the cosmic structures and of the IGM as well as the efficiency for its reionization. Compared to previous studies, we do not restrict our discussion to the global metallicity but we specifically follow several chemical elements (D, $^{4}$He, C, N, O, Si, S, Fe, Zn). We further discuss the implications of the observation of the extremely metal poor star ([Fe/H] = -5.3), HE 0107-5240. This star can severely constrain such models and/or models of stellar nucleosynthesis, particularly those for extremely massive stars which evolve through pair instability. We find that a bimodal (or top-heavy) IMF (40 - 100 M$_\\odot$), for example as discussed by \\citet{ciardi:03}, is preferred over a scenario which includes a component with very massive stars ($\\ga$ 100 M$_\\odot$) which are often assumed to be responsible for the early stage of reionization. The best model is well suited to account for C and Zn abundances, as well as Fe, O, Si and S in the DLAs, in the metal-poor IGM, and in the oldest halo stars. Furthermore, in order to reionize the universe at $z\\sim 17$, we require an escape fraction for ionizing photons of only about 5 \\%. In contrast, we find that the massive starburst model, relying mostly on pair instability supernovae (PISNae) from stars in the range 140--260 M$_{\\odot}$, requires a lower SFR intensity due to the overproduction of metals and as a consequence, the UV flux is weaker and requires a higher escape fraction of ionizing photons ($\\sim$ 30 \\%) to achieve early reionization. Note that \\citet{venkatesan:03b} have also found a low ionizing efficiency for very massive stars. Nevertheless, the overproduction of some $\\alpha$--elements (specifically S and Si) seems unavoidable in these models. Moreover due to the specific yields of these PISNae, the observations of the metal-poor halo stars are not reproduced. We also consider a model with an extremely massive component, i.e., with a stellar mass range $\\ge$ 270 M$_{\\odot}$. Assuming that these massive stars collapse entirely into black holes in this mode, chemical evolution and reionization are de-correlated. We find that the ionizing flux from these very massive stars can very easily reionize the Universe at $z\\sim 17$. However we also find that the resulting chemical evolution in this model is exactly the same as in a standard star formation model, and the high redshift abundances are not reproduced. The outline of the paper is as follows: in section 2, we describe the chemical evolution models in detail, and our method for calculating the chemical enrichment inside the structures and the IGM. In section 3, we describe our computation of the flux of ionizing photons. Our results are collected in section 4 and our conclusions are summarized in section 5. ", "conclusions": "It is likely that ionizing photons from stars, as opposed to say accreting black holes (in effect quasars or miniquasars) are responsible for the early epoch of reionization at $z\\sim 15-20$ inferred from the WMAP results \\citep[cf.][]{dijkstra:04a}. However there are severe limitations on massive stars if these formed with a ``normal'' IMF. The hypothesis of a generation of very massive metal-free stars helps supply the ionization requirements \\citep*{cen:03a,wyithe:03}. However, we show here that chemical evolution constraints essentially rule out such an interpretation. Rather, we describe how a standard massive IMF ($40$-$100\\ \\mathrm{M}_{\\odot}$), for example as discussed by \\citet{ciardi:03}, can reionize the early intergalactic medium (see also a discussion of an even less massive IMF by \\citealt*{venkatesan:03b}). Our approach is to use the cosmic star formation rate history as our primary guide to early star formation rates, and work with integrated quantities in the context of hierarchical structure formation to yield the ionizing photon flux history for alternative cosmic star formation models which are capable of reionizing the early intergalactic medium. We have developed models which include an early burst of massive stars combined with standard star formation. We computed the stellar ionizing flux of photons and we tracked the nucleosynthetic yields for several elements\\,: D, $^{4}$He, C, N, O, Si, S, Fe, Zn as a function of redshift, both in the IGM and in the ISM of the growing structures. We compared the results of these models with the observed abundances in the Lyman $\\alpha$ forest and in DLAs. We also considered possible constraints associated with the observations of the two extremely metal-poor stars HE 0107-5240 and CS22949-037. We have shown that a bimodal (or top-heavy) IMF (40 - 100 M$_\\odot$) best satifies all constraints applied. In contrast, models with an extreme mass range ($\\ga$ 100 M$_\\odot$) often assumed to be responsible for the early stages of reionization do less well. As motivated by the numerical simulations of first star formation \\citep*{abel:02,bromm:02}, a mode of even more extreme stellar masses in the range ($\\ge$ 270 M$_{\\odot}$) has also been considered. As all massive stars collapse entirely into black holes in this mode, the chemical evolution and the reionization are de-correlated, as already mentioned by \\citet{tumlinson:04}. The ionizing flux from these very massive stars can easily reionize the Universe at $z\\sim 17$. However the chemical evolution in this case is exactly the same as in the standard star formation model, and the high redshift abundances are not reproduced. Consequently, the suggestion \\citep*{bromm:03} that such Population III stars were the precursors of the extreme metal-poor halo stars is untenable. There is no evidence, nor any need, for a hypothesised primordial population of very massive % stars in order to account for the chemical abundances of extremely metal-poor halo stars or of the intergalactic medium. The combined population of early-forming, normal (0.1 - 100 M$_\\odot$) and massive (40 - 100 M$_\\odot$) stars can simultaneously explain the cosmic chemical evolution and the observations of extremely metal-poor halo stars and also account for early cosmological reionization. We have shown that the initial massive starburst, which originally was introduced to reionize the early Universe, produces rapid initial metal pollution. The existence of old, C-rich halo stars with high [O/Fe] and [C/Fe] ratios is predicted as a consequence of these massive stars. The recently observed abundances in the oldest halo stars could trace this very specific stellar population. We have also found that the D abundance is strongly coupled to the gas fraction in the structures, with the implication that local D measurements are a non-robust cosmological probe. In addition, there is some non-primordial contribution to the He abundance even in metal-poor galaxies, that is however within the observed range. Our suggestion is far from the whole story. There must be late ionization input by harder photons than produced by OB stars to account for the \\ion{He}{2} abundance at $z\\sim 2-4,$ for example associated with the quasar population \\citep*{hui:03} or Pop III stars \\citep*{bromm:01a,venkatesan:03a}. As structure builds up, the ionizing photon escape fraction will surely decrease, and the intergalactic medium is likely to recombine before $z\\sim 6,$ at which time the neutral fraction is constrained from the onset of the Gunn-Peterson effect to be about 0.001\\%. \\citet{cen:03b} has suggested that the Universe becomes neutral again at $z \\sim 13$ and is reionized for the second time at $z\\sim 6$, not necessary by the same stellar population. However our calculations suggests that normal star formation is likely to suffice at this epoch \\citep{gnedin:04} in order to account for the observed ionization level." }, "0405/astro-ph0405449_arXiv.txt": { "abstract": "Optical and near-infrared observations of the gamma-ray burst GRB\\,031203, at $z = 0.1055$, are reported. A very faint afterglow is detected superimposed to the host galaxy in our first infrared $JHK$ observations, carried out $\\sim 9$~hours after the burst. Subsequently, a rebrightening is detected in all bands, peaking in the $R$ band about 18 rest-frame days after the burst. The rebrightening closely resembles the light curve of a supernova like SN\\,1998bw, assuming that the GRB and the SN went off almost simultaneously, but with a somewhat slower evolution. Spectra taken close to the maximum of the rebrightening show extremely broad features as in SN\\,1998bw. The determination of the absolute magnitude of this SN (SN\\,2003lw) is difficult owing to the large and uncertain extinction, but likely this event was brighter than SN\\,1998bw by $0.5$~mag in the $VRI$ bands, reaching an absolute magnitude $M_V = -19.75 \\pm 0.15$. ", "introduction": "In recent years, extensive optical and near-infrared (NIR) follow-up of gamma-ray bursts (GRBs) has revealed a physical connection between a significant fraction of long-duration GRBs and core-collapse supernov\\ae{} (SNe). First, the bright SN\\,1998bw was discovered spatially and temporally coincident with GRB\\,980425 \\citep{Ga98,Ku98}. However, GRB\\,980425 was rather different from classical, cosmological GRBs, being severely underenergetic and lacking an optical afterglow. Then, SN\\,2003dh was detected in the afterglow of GRB\\,030329 \\citep{St03,Hj03}. Both SNe showed broad bumps in their spectra, indicating very large expansion velocities (up to $30\\,000$~km/s), and were extremely bright. These highly-energetic SNe are often named hypernov\\ae{} \\citep[e.g.][]{Iw98}. Last, bumps discovered in the light curves of several afterglows, peaking $\\sim 20$~days after the GRB, have been interpreted as due to the emerging of SNe out of the afterglow light, based on their brightness, temporal evolution and colors \\citep[e.g.][]{Bl99,Ga03}. The bumps resemble the light curve of SN\\,1998bw, with a certain scatter in the brightness and rise time \\citep[e.g.][]{Ze04}. Spectroscopic confirmation that the bump of GRB\\,021211 has a SN spectrum \\citep[SN\\,2002lt;][]{DV03} supports this conclusion. These observations indicate that the GRB/SN association is common. GRB\\,031203 was discovered by the INTEGRAL satellite on 2003 Dec 3.91769 UT \\citep{Go03}, with a duration of $\\sim 30$~s and a peak flux of $1.3 \\times 10^{-7}$~erg~cm$^{-2}$~s$^{-1}$ \\citep[$20-200$~keV;][]{Me03a}. The precise and fast dissemination of the GRB coordinates by the INTEGRAL burst alert system \\citep{Me03b} allowed an effective search for the afterglow. We also immediately activated our ToO program at ESO, starting NIR observations at the NTT 7~hours after the GRB \\citep{Ze03}. The X-ray and radio afterglows were soon discovered \\citep{Sa03,Fr03}. A compact galaxy, located at a consistent position, was proposed to be the GRB host galaxy by \\citet{Pr03}. The redshift was $z = 0.1055 \\pm 0.0001$ \\citep{Pr03,Pr04}, making GRB\\,031203 the second closest burst after GRB\\,980425 at $z = 0.0085$ \\citep{Ga98}. \\citet{Va04} discovered a scattered expanding X-ray halo due to the reflection of the burst and/or early afterglow light from Galactic dust grains. This allowed an (indirect) measurement of the X-ray flux at the earliest stages after the burst onset. Given the low redshift of this event, the isotropic-equivalent burst energy is extremely low% \\footnote{We adopt a cosmology with $H_0 = 71$~km~s$^{-1}$~Mpc$^{-1}$, $\\ped{\\Omega}{m} = 0.27$, $\\Omega_\\Lambda = 0.73$ (WMAP results). At $z = 0.1055$ the luminosity distance is $D =477$~Mpc and the distance modulus is $\\mu = 38.42$~mag.}, $\\ped{E}{iso} \\sim 3 \\times 10^{49}$~erg \\citep[20-2000~keV;][]{Wa04,Pr04}, well below the standard reservoir $\\sim 2 \\times 10^{51}$~erg of normal GRBs \\citep{Fr01,Bl03}. Only GRB\\,980425 \\citep{Ga98} and XRF\\,020903 \\citep{Sa04} were less energetic. Based on photometric monitoring of the host galaxy, several groups have reported evidence for a SN associated with GRB\\,031203 \\citep{Be04,Th04,Co04,Ga04}. After the ultimate confirmation, coming from spectroscopic observations and reported by our group \\citep{Ta04}, the IAU named this event SN\\,2003lw. ", "conclusions": "\\begin{figure}\\centering \\includegraphics[width=\\columnwidth]{lc.ps} \\caption{Optical and NIR light curves of GRB\\,031203 (dots). Error bars indicate the amount of relative errors only (Tab.~\\ref{tb:phot}). The solid curves show the evolution of SN\\,1998bw \\citep{Ga98,Mc00}, rescaled at $z = 0.1055$, stretched by a factor 1.1, extinguished with $E_{B-V} = 1.1$ and brightened by 0.5~mag. Dashed lines indicate the host galaxy contribution. Vertical lines mark the epochs of our spectra.\\label{fg:lc}} \\end{figure} \\begin{figure}\\centering \\includegraphics[width=\\columnwidth,keepaspectratio]{SED_nuFnu.ps} \\caption{Spectral energy distribution of the afterglow of GRB\\,031203 on 2003 Dec.~4.3 UT (9 hours after the trigger). The NIR values are calculated from our data by subtracting the host contribution and assuming $E_{B-V} = 1.1$. The NIR spectral index is $\\beta = 2.36 \\pm 0.02$ ($F_\\nu \\propto \\nu^{-\\beta}$). The X-ray spectrum is from \\citet{Wa04} (reported at $t = 9$~h using the X-ray decay slope), who find $\\ped{N}{H} = 8.8 \\times 10^{21}$~cm$^{-2}$ and $\\ped{\\beta}{X} = 0.90 \\pm 0.05$.\\label{fg:SED}} \\end{figure} In Fig.~\\ref{fg:lc} we show the light curves of GRB\\,031203. Early-time NIR photometry shows a dimming in all bands between the first and second night after the GRB. This is confirmed by PSF-matched image subtraction. We believe that we have seen the NIR afterglow of GRB\\,031203. The magnitudes are $J = 20.60 \\pm 0.09$, $H = 19.05 \\pm 0.07$, $K = 17.56 \\pm 0.05$ (9~hours after the GRB), obtained by subtracting the host contribution. \\citet{Co04} have $I$-band observations at similar epochs, and do not report evidence for variability. However, extrapolation to the visible region yields $I \\sim 23.4$, quite a faint value when compared to the host luminosity $I \\approx 19.4$. Little contribution from the afterglow is seen in our measurement of Dec.~5, implying a quick decay between the two nights ($F(t) \\propto t^{-\\alpha}$, with $\\alpha \\gtrsim 2$). However, there is no variation between the two $K$-band observations of the first night (separated by 2.6~h), suggesting a break in the light curve or a bumpy behaviour. In Fig.~\\ref{fg:SED}, we compare the spectrum in the NIR and X-ray regions \\citep{Wa04}. A discontinuity is apparent, indicating a different origin for the emission in the two bands. The X-ray component has a much harder spectrum, and a slower decay ($\\alpha = 0.55 \\pm 0.05$). Interestingly, \\citet{Wa04} infer a fast decay of the early-time X-ray afterglow, consistent with our NIR value ($\\alpha \\gtrsim 1.7$). In the standard model of afterglows \\citep[e.g.][]{SPN98}, a fast decay is consistent with a soft spectrum bluewards of the peak frequency. A few days after the GRB, a rebrightening is apparent in all optical/NIR bands. The rebrightening amounts to $\\approx 30\\%$ of the total flux, and is coincident with the center of the host galaxy to within $0.1\\arcsec$ ($\\approx 200$~pc at $z = 0.1055$). For comparison, we show in Fig.~\\ref{fg:lc} the $VRI$ light curves of SN\\,1998bw \\citep{Ga98,Mc00}, placed at $z = 0.1055$ and dereddened with $E_{B-V} = 1.1$ (see below). Interpolation of the $UBVRI$ data was performed in order to estimate the fluxes of SN\\,1998bw at the frequencies corresponding to the observed bands. Even after correcting for cosmological time dilation, the light curve of SN\\,2003lw is broader than that of SN\\,1998bw, and requires an additional stretching factor of $\\approx 0.9$ to match the $R$ and $I$ bands. Near the peak, the light curve is rather flat, resembling the hypernova SN\\,1997ef \\citep{Iw00} more than SN\\,1998bw. The $R$-band maximum is reached on approximately 2003 Dec.~24 ($\\sim 18$ comoving days after the GRB). We note that the details of the light curve shape are sensitive to the removal of the host contribution. This may explain the different finding of \\citet{Th04}, who need no stretch, and \\citet{Co04}, who find a longer rise. Assuming a light curve shape similar to SN\\,1998bw, which had a rise time of 16~days in the $V$ band, our data suggest an explosion time nearly simultaneous with the GRB. However, given that SN\\,2003lw was not strictly identical to SN\\,1998bw, and as we lack optical data in the days immediately following the GRB, a lag of a few days cannot be ruled out. Type-Ic SNe usually reach $V$-band maximum in $\\sim 12$-20 days, the brightest events showing a slower evolution \\citep[see e.g. Fig.~2 of][]{Maz02}. A precise determination of the absolute magnitude of the SN is made difficult by the uncertain, and significant, extinction. C04 and \\citet{Pr04} constrain the average combined Galactic and host extinction to be $E_{B-V} \\approx 1.1$ based on the Balmer ratios of the host galaxy. Given the good spatial coincidence of the SN with the center of the host, such value is likely a good estimate for the SN extinction. We also adopt a Galactic extinction law \\citep{Ca89} with $R_V = 3.1$. With the assumed reddening, SN\\,2003lw appears brighter than SN\\,1998bw by 0.5~mag in the $V$, $R$, and $I$ bands. The absolute magnitudes of SN\\,2003lw are hence $M_V = -19.75\\pm0.15$, $M_R = -19.9\\pm0.08$, and $M_I = -19.80\\pm0.12$. \\citet{Th04}, using $I$-band data, also found that SN\\,2003lw was brighter than SN\\,1998bw by $\\sim 0.55$~mag, in full agreement with our result. \\citet{Co04}, however, found a comparable luminosity for the two SNe; this discrepancy is entirely due to the lower extinction they assume. Fig.~\\ref{fg:spec} shows the spectra of the rebrightening on 2003 Dec.~20 and Dec.~30 (14 and 23 rest-frame days after the GRB), after subtracting the spectrum taken on 2004 Mar.~1 (81 rest-frame days after the GRB). This assumes that the latter spectrum contains only a negligible contribution from the SN, which is confirmed by the photometry (Fig.~\\ref{fg:lc}). The spectra of SN\\,2003lw are remarkably similar to those of SN\\,1998bw obtained at comparable epochs \\citep[shown as dotted lines in Fig.~\\ref{fg:spec}; from][]{Pa01}. Both SNe show very broad absorption features, indicating large expansion velocities. Thus we tentatively classify SN\\,2003lw as a hypernova. The main absorptions are identified in Fig.~\\ref{fg:spec} as in SN\\,1998bw, following \\citet{Iw98}. The velocity of the Si\\,II line in SN\\,2003lw is apparently smaller than in SN\\,1998bw. The broad peaks near 5300~\\AA{} and 6600~\\AA{} are probably the emission components of P-Cygni profiles due to the blending of several lines. There is evolution between the two epochs: the bluer bump is observed at longer wavelengths in the second spectrum, and is slighty narrower. Moreover, the shape of the redder peak is different in the two epochs. Both peaks appear at redder wavelengths than in SN\\,1998bw. Detailed modeling of the spectra will be presented elsewhere (Mazzali et al. 2004, in preparation).\\bigskip \\begin{figure}\\centering \\includegraphics[width=\\columnwidth,keepaspectratio]{spec.ps} \\caption{Spectra of SN\\,2003lw, taken on 2003 Dec.~20 and Dec.~30 (solid lines), smoothed with a boxcar filter 250~\\AA{} wide. Dotted lines show the spectra of SN\\,1998bw \\citep[from][]{Pa01}, taken on 1998 May~9 and May~19 (13.5 and 23.5~days after the GRB, or 2 days before and 7~days after the $V$-band maximum), extinguished with $E_{B-V} = 1.1$ and a Galactic exinction law \\citep{Ca89}. The spectra of SN\\,1998bw were vertically displaced for presentation purpose.\\label{fg:spec}} \\end{figure} By modeling the X-ray dust echo, \\citet{Wa04} concluded that GRB\\,031203 was an X-ray flash (XRF); however, the prompt emission data do not confirm this hypotesis (Sazonov, Lutovinov \\& Sunyaev 2004, in preparation). This event, like SN\\,1998bw \\citep{Pi00}, seems therefore to violate the correlation between the isotropic-equivalent gamma-ray energy $\\ped{E}{iso}$ and the peak spectral energy $\\ped{E}{p}$ \\citep{Am02,La03}. In fact, assuming $\\ped{E}{iso} \\sim 1.5 \\times 10^{50}$~erg \\citep[1-10000~keV;][]{Wa04}, the \\citet{Am02} relation would imply $\\ped{E}{p} \\sim 10$~keV, a value indicating an XRF nature for GRB\\,031203. This is in contrast with INTEGRAL data. Of course, this issue will be be settled only after a thorough analysis of the prompt emission spectra. The afterglow of GRB\\,031203 was very weak, the faintest ever detected in the optical/NIR. Extrapolation in the $R$ band yields a luminosity $\\sim 200$~times fainter than the dimmest afterglow discovered so far \\citep[GRB\\,021211:][]{Fo03,Pa03}. The detection of the SN optical light implies that the reason of such faintness was not an extreme dust obscuration. Also given the low redshift of the event, this example shows that some optical afterglows may escape detection just because they are faint \\citep[e.g.][]{Fy01,La02,DP03}. GRB\\,031203, together with GRB\\,980425 at $z = 0.085$, was a very dim event, perhaps a jet observed far from its axis \\citep[e.g.][]{Ma02,Ya03}. Being so faint, they would have been likely missed at cosmological distances. Since the volume they sample is much smaller than that probed by classical, distant GRBs with $\\langle z \\rangle \\approx 1$, the rate of these events could be much larger. As noted by \\citet{Th04}, this would increase the detection rate for the {\\it Swift} satellite \\citep{Ge04}. More rapid and efficient observations, also soon feasible thanks to {\\it Swift}, will allow a detailed study of this largely unexplored class of events. GRB\\,031203 was quite similar to GRB\\,980425, even if overall more powerful. Both events consisted in a single, underenergetic pulse. Their afterglows were very faint or absent in the optical, and showed a very slow decline in the X-ray \\citep{Pi00,Wa04}. Last, they were both accompained by a powerful hypernova." }, "0405/astro-ph0405163_arXiv.txt": { "abstract": "We present a new mechanism for Type Ia supernova explosions in massive white dwarfs. The proposed scenario follows from relaxing the assumption of symmetry in the model and involves a detonation created in an unconfined environment. The explosion begins with an essentially central ignition of stellar material initiating a deflagration. This deflagration results in the formation of a buoyantly-driven bubble of hot material that reaches the stellar surface at supersonic speeds. The bubble breakout forms a strong pressure wave that laterally accelerates fuel-rich outer stellar layers. This material, confined by gravity to the white dwarf, races along the stellar surface and is focused at the location opposite to the point of the bubble breakout. These streams of nuclear fuel carry enough mass and energy to trigger a detonation just above the stellar surface. The flow conditions at that moment support a detonation that will incinerate the white dwarf and result in an energetic explosion. The stellar expansion following the deflagration redistributes stellar mass in a way that ensures production of intermediate mass and iron group elements consistent with observations. The ejecta will have a strongly layered structure with a mild amount of asymmetry following from the early deflagration phase. This asymmetry, combined with the amount of stellar expansion determined by details of the evolution (principally the energetics of deflagration, timing of detonation, and structure of the progenitor), can be expected to create a family of mildly diverse Type Ia supernova explosions. ", "introduction": "Type Ia supernovae are one class of luminous stellar explosions. These are the predominant explosive events in old stellar environments such as elliptical galaxies. The ejecta of these objects are rich in intermediate mass and iron group elements. Explaining the nature of these objects is therefore critical for understanding galactic chemical evolution \\citep{truran+71}. These supernovae also are the key component of one method used to determine the history of the Universe and probe the origin of dark energy \\citep{sandage+93,perlmutter+99,tonry+03,knop+03}. The work presented in this {\\em Letter} builds on many previous observational and theoretical contributions to the field of Type Ia SNe, and on advances in fluid dynamics, nuclear physics, and computational science. Current ideas about the Type Ia supernova explosion mechanism follow from the original work of Arnett, Nomoto, and Khokhlov \\citep{arnett69,nomoto+76,khokhlov91-dd}, who pioneered deflagrating and detonating massive white dwarfs originally proposed by \\cite{hoyle+60} as the core component of Type Ia supernovae. Despite decades of effort, these events remain an unsolved mystery. Proposed explosion scenarios include white dwarf detonations \\citep{arnett69,nomoto82}, coalescing pairs of white dwarfs \\citep{webbink84,iben+84}, deflagrations or delayed detonations of massive white dwarfs \\citep{nomoto+84,khokhlov91-dd}, and collapse in a strong gravitational field \\citep{wilson+04}. None of these scenarios accounts for all the observed features of Type Ia supernovae. Some models produce energetic explosions but fail to explain the observed ejecta compositions, while others successfully produce the observed chemically stratified ejecta but require including ad hoc physics. Here we report the results of multi-dimensional hydrodynamical simulations of the long-term evolution of a massive white dwarf following an essentially central ignition. The initial evolution results in a deflagration front and formation of a Rayleigh-Taylor unstable buoyancy-driven bubble that is accelerated to supersonic speeds on its way to the stellar surface. We follow the evolution beyond bubble breakout and observe formation of gravitationally confined flow across the surface of the star. The flow is focused at the point opposite the breakout location, where the matter is compressed and heated, igniting a detonation wave just above the stellar surface. We find that the star expands substantially during the evolution. This expansion produces a density distribution that, when overrun by the detonation wave, can be expected to result in ejecta that are strongly layered and rich in intermediate mass elements, as is typical of Type Ia supernovae. The modest asymmetry in the expansion of the star can also be expected to produce a mild amount of global asymmetry in the explosion. ", "conclusions": "We presented a gravitationally confined detonation (GCD) mechanism for Type Ia supernova explosions. The basic components of the proposed scenario are a rising deflagrating bubble expelling a small amount of stellar matter, the associated stellar expansion caused by the shallower potential well, the flood of stellar material across the surface following the bubble breakout, and a detonation in an gravitationally confined environment. In the GCD scenario, not imposing any asymmetries in the initial conditions is of paramount importance. A deflagration front is born very close to the white dwarf center. Such an essentially central ignition is more probable than the idealized conditions adopted in standard deflagration or delayed detonation models, primarily because the central region of the star is convective. This type of ignition results in a rising deflagrating bubble accelerated by buoyancy to supersonic speeds on its way to the stellar surface. The transonic phase of the bubble's rise is accompanied by the formation of a bow shock ahead of the bubble that compresses and heats the nuclear fuel. Our attempts to associate this region with a possible transition to detonation failed. We discovered, however, that the flood of the expelled surface layers that follows the bubble breakout remains confined to low radii above the star. This flood races around the star and is ultimately focused into a hot, compressed, high density region located just above the stellar surface. Conditions in this region satisfy the criteria necessary for a detonation. One important aspect of the GCD mechanism is that stellar expansion is a natural consequence of the essentially central ignition of a deflagration. The flame releases energy and displaces mass, softening the gravitational potential well leading to expansion of the star. The expansion will slow down on a time scale comparable to the sound crossing time of the white dwarf as the star tries to reach a new equilibrium. Therefore, we expect that at still later times, if not for the fact that the detonation will completely disrupt the white dwarf, the initial expansion would be followed by contraction of stellar material and the star would oscillate. Because of this pre-expansion, the detonation front born above the stellar surface will encounter densities similar to those found in models where pre-expansion results from a centrally ignited large-scale deflagration~\\citep{reinecke+02,gamezo+03}. The estimate of nucleosynthetic yield for intermediate mass (iron peak) elements is a lower (upper) limit in view of the fact that the stellar expansion continues after the moment of the detonation. Determining the final yields, however, requires simulating the detonation. The actual conditions across the detonation wave, particularly the amount of compression, will influence the results. Also, the yields can be modified by delaying or by accelerating the ignition of the detonation to create a family of Type Ia supernovae with diverse spectral characteristics. The exact timing of the detonation depends on several factors. The structure of the progenitor influences energy release by the deflagration. It also affects the strength and mass of the stellar layers being pushed around the star. The radius of the progenitor regulates the time required by the streaming matter to reach the confluence point. All these factors determine the time available for stellar pre-expansion and will be a source of diversity in GCD models. Some properties of the proposed model are, however, largely independent of the precise details of the ignition of the detonation or stellar progenitor. The explosion will be powerful. All the stellar fuel will be consumed by the detonation. Despite the perturbation introduced by the deflagration, the star will retain most of its radially symmetric stratification. Therefore, the subsequent explosion will display characteristics well-known from one-dimensional investigations \\citep{nomoto+84,hoeflich+98}. In particular, we expect the distribution of nucleosynthetic products in velocity space to agree with the observed layered structure of Type Ia supernova ejecta. The model also naturally admits certain asymmetries. The deflagration consumes only about 5\\% of the stellar mass. We expect a similar level of variation in the resulting spectra and luminosities, in agreement with degree of diversity present in the observations \\citep{li+01}. In addition, because the stellar shape is distorted by the rising bubble and the formation of the detonation on one side of the star, we expect a noticeable asymmetry in the stellar ejecta. These orientation effects might be responsible for peculiar events such as SN 1991T \\citep{filippenko97}. Because the gross properties of observed Type Ia supernovae can be accounted for in the GCD model, detailed observations will be required to verify the proposed mechanism. One possibility is to obtain information about the degree of asymmetry in these explosions, using precision polarimetric measurements. In this short communication, we have not presented simulations that support all of our predictions. Rather, we have presented a logical sequence of events that naturally leads to a Type Ia supernova explosion. Several of these steps need to be carefully studied. In particular, the proposed detonation mechanism bears many similarities to the process of confined fusion studied in terrestrial laboratory experiments, which is notoriously difficult and prone to instabilities. For this reason, the early stages of the detonation should be studied very carefully. The expected computational demands and challenges are severe and clearly approach limits of feasibility. Despite the fact that such a study lies in the future, we are confident in the basic components of the proposed GCD mechanism." }, "0405/astro-ph0405213_arXiv.txt": { "abstract": "We present the results of \\xmm\\ observations of three high-redshift powerful radio galaxies 3C\\,184, 3C\\,292 and 3C\\,322. Although none of the sources lies in as rich an X-ray-emitting environment as is seen for some powerful radio galaxies at low redshift, the environments provide sufficient pressure to confine the radio lobes. The weak gas emission is particularly interesting for 3C\\,184, where a gravitational arc is seen, suggesting the presence of a massive cluster. Here {\\it Chandra\\/} data complement the \\xmm\\ measurements by spatially separating X-rays from the extended atmosphere, the nucleus and the small-scale radio source. For 3C\\,292 the X-ray-emitting gas has a temperature of $\\sim$2 keV and luminosity of 6.5$\\times10^{43}$ erg s$^{-1}$, characteristic of a poor cluster. In all three cases, structures where the magnetic-field strength can be estimated through combining measurements of radio-synchrotron and inverse-Compton-X-ray emission, are consistent with being in a state of minimum total energy. 3C\\,184 and 3C\\,292 (and possibly 3C\\,322) have a heavily absorbed component of nuclear emission of $N_{\\rm H} \\sim $ few 10$^{23}$ cm$^{-2}$. ", "introduction": "Powerful radio galaxies are visible to high redshift and hence can be used as pointers to large-scale structures. It has been hypothesised that powerful radio galaxies may represent a means of discovering high-redshift clusters of galaxies (e.g., Le F\\`evre et al. 1996; Fabian et al. 2001) through their X-ray emission, with the possibility of placing constraints on cosmological parameters. Analysis of pointed ROSAT observations of $z>0.5$ radio galaxies and quasars in the 3CRR catalogue (Laing, Riley \\& Longair 1983) gave support to this hypothesis, resulting in claims of the detection of hot ($\\sim$ 5 keV), luminous ($L_{\\rm X}>10^{44} h_{50}^{-2}$ ergs s$^{-1}$) gas around a number of sources (Crawford et al. 1993, Worrall et al. 1994; Crawford \\& Fabian 1995, 1996; Hardcastle et al. 1998a; Hardcastle \\& Worrall 1999; Crawford et al. 1999). However, because of the limited sensitivity and spectral resolution of the instruments on-board ROSAT, there was no spectral confirmation of these results, and the detected emission was spatially resolved in relatively few cases (see Hardcastle \\& Worrall 2000). With the advent of sensitive X-ray telescopes, it is now possible to detect and study the environment of such galaxies as a function of cosmic time. Several physical mechanisms are responsible for the X-ray emission from radio galaxies. Thermal emission from a hot atmosphere surrounding a radio galaxy gives information about the large scale gas distribution and the interaction between the expanding radio source and the external medium. Unresolved emission from the active galactic nucleus (AGN) may consist of thermal or non-thermal components, and probes the physical conditions near the central engine. The radio jets, lobes and hot spots can also emit detectable X-rays, by the inverse Compton (IC) and/or synchrotron processes. IC X-ray measurements can be combined with radio data to give information about particle acceleration and magnetic field strength. Powerful radio galaxies at intermediate redshift are found in rich, cluster-like environments (Hill \\& Lilly 1991) and some (but not all) low-redshift high-power objects such as Cygnus A (e.g. Smith et al. 2002; Markevitch, Sarazin \\& Vikhlinin 1999) and Hydra A (McNamara et al. 2000; Nulsen et al. 2002) also lie in rich clusters which have been studied in detail. The X-ray/radio properties of FRII galaxies of low/intermediate redshift have been studied with ROSAT (e.g., Carilli, Perley \\& Harris 1994; Leahy \\& Gizani 1999; Hardcastle \\& Worrall 2000). Hardcastle \\& Worrall (2000) estimated the gas pressure in the X-ray external medium and compared it to the internal pressure of the radio-emitting plasma assuming the condition of minimum total energy in the magnetic field and radiating electrons. They concluded that the majority of the sources might be underpressured with respect to the external medium, and discussed additional contributions to the internal pressure that would avoid collapse of the radio lobes. \\begin{table*} \\caption{Source specific details} \\label{tab:sources} \\begin{center} \\begin{tabular}{rcccccc} \\hline Source\t& $\\alpha$(h m s)& $\\delta$ (o \\arcmin ~\\arcsec) & redshift & $N_{\\rm H}^{(3)}$ & Flux density at 178 MHz$^{(4)}$ & scale \\\\ & (J2000)\t& (J2000) \t\t\t &\t & 10$^{20}$ (cm$^{-2}$) & (Jy) & kpc/arcmin \\\\ \\hline 3C 184\t&07 39 24.30\t& +70 23 10.7\t& 0.994$^{(1)}$\t& 3.45 & 13.2 & 480 \\\\ 3C 292\t&13 50 41.95\t& +64 29 35.4\t& 0.710$^{(1)}$ & 2.17 & 10.1 & 431 \\\\ 3C 322\t&15 35 01.16 \t& +55 36 51.4\t& 1.681$^{(2)}$ & 1.31 & 10.1 & 508 \\\\ \\hline \\end{tabular} \\vskip 5pt \\end{center} \\begin{minipage}{13. cm} (1) From Nilsson \\cite{Nilsson98}- (2) from Strom et al. \\cite{strometal}. (3) $N_{\\rm H}$ values are from Dickey \\& Lockman \\cite{nh} - (4) Flux density at 178 MHz is from the 3CRR catalog - Laing, Riley \\& Longair \\cite{LRL83}. \\end{minipage} \\end{table*} The new generation X-ray satellites, \\xmm\\ and {\\it Chandra\\/}, are well suited for studies of high-redshift radio galaxies, the first for its unrivalled sensitivity, the second for its excellent spatial resolution. {\\em Chandra} observations of radio galaxies and quasars at redshift $z>0.5$ have found relatively few sources in the rich gaseous environments associated with massive clusters. Most claims are based on the flux of the extended emission (e.g. Hardcastle et al. 2002; Crawford \\& Fabian 2003), with only a few confirmed by spectral measurements (3C\\,220.1, $kT = 5$ keV --- Worrall et al. 2001; 3C\\,294, $kT = 3.5$ keV --- Fabian et al. 2003, although a large non-thermal contribution to the emission is found to be possible.). The detected atmospheres of other sources are more typical of a group or poor cluster (e.g. Crawford \\& Fabian 2003, Donahue et al. 2003). Although the very presence of edge-brightened radio lobes points to the existence of {\\em some} gas in order that the lobes should be confined, a particularly rich environment is not required, and the {\\it Chandra} observations seem to point in this direction. It is interesting to note that for {\\em Chandra}-observed radio galaxies and quasars, the diffuse X-ray emission is usually elongated in the direction of the galaxy radio-lobe axis (e.g. Donahue et al. 2003; Carilli et al. 2002, Hardcastle et al. 2002). This has been interpreted as showing that some, if not most, of the X-ray extended emission is not thermal but arises from synchrotron or IC process (see also Fabian et al. 2003). The dominant inverse-Compton-scattered radiation field is expected to be the Cosmic Microwave Background (CMB), but on small scales this may be supplemented by photons from the nucleus and radio source (Brunetti 2000). In this paper, we present a study of high-redshift radio galaxies based on \\xmm\\ observations. The large collecting area of \\xmm\\ provides an unprecedented opportunity for detecting relatively poor clusters around radio galaxies at $z \\sim 1$, as well as emission associated with extended radio components. Measurements of the nuclear X-ray emission help to test models which seek to unify high-redshift radio galaxies and quasars based on orientation, assuming the presence of a central absorbing torus of gas and dust. Thus we might expect at least moderate intrinsic absorption of the nuclear emission from radio galaxies, although in {\\it Chandra\\/} observations to date the nucleus more commonly shows low intrinsic absorption, suggesting that emission from an inner jet is dominant (Worrall et al. 2001). Here we discuss the X-ray results for three sources: 3C\\,184 ($z$=0.994), 3C\\,292 ($z$=0.710) and 3C\\,322 ($z$=1.681), see Table \\ref{tab:sources}. The sources are all radio galaxies at $ z >0.5$ selected from the 3CRR catalogue. Their selection was made on the basis of \\xmm\\ observing constraints, i.e., no nearby bright X-ray sources which might enhance the background from stray light, no nearby bright stars, and high \\xmm\\ visibility. They were also chosen to be sources with no ROSAT, {\\it Chandra}, or \\xmm\\ observations made or planned at the time they were proposed, and to enhance the sensitivity at low energies we required the Galactic $N_H$ to be less than $5 \\times 10^{20}$ cm$^{-2}$. Thus, since they were not selected for intrisic source characteristics, the X-ray results should be typical of the 3CRR radio-galaxy population at high redshift. The two at lower redshift belong to a larger sample of 33 radio galaxies and quasars that will be observed with {\\it Spitzer} in a guaranteed-time programme. Throughout the paper we use a cosmology with $H_{\\rm 0}$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m}$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7. If not otherwise stated, errors are quoted at 1$\\sigma$ confidence level. ", "conclusions": "\\subsection{The origin of the X-ray emission of 3C 184} The combined spectro-imaging analysis of 3C\\,184 indicates that half the X-ray emission is produced by an absorbed component which is associated with the AGN nucleus. We see also non-thermal emission from a soft component which is likely to be radio related, at least in part. Other powerful radio galaxies (e.g. Hardcastle et al. 2002; Donahue et al. 2003) show soft X-ray emission associated with the position of the outer radio lobes. By comparing the \\xmm\\ observation with the 20 ks Chandra observation, the emission from the lobes is separated from the nuclear emission. Only 5 {\\em Chandra} counts correspond to the region of the radio lobes, not sufficient for a spectral analysis. However, we estimated an X-ray flux density of 0.2$\\pm0.1$ nJy at 1 keV, which is in excellent agreement with the flux predicted by a model where the source of scattered photons is a combination of CMB and IR photons from the nucleus. This allows us to conclude that the radio-galaxy lobe is very close to equipartition between the magnetic field and relativistic electrons. We have marginal evidence, from the spatial analysis, for a third component corresponding to a gaseous environment. The luminosity of the extended component estimated from the $\\beta$-model fitting (with $\\beta$ fixed at 0.66) is 5.9$\\times10^{43}$ erg s$^{-1}$ within the detection radius of 88 arcsec, which rises to $L_{\\rm X}( 0.4$ (Vikhlinin et al. 2002). Is this medium rich enough to confine the expansion of the radio galaxy? The average external pressure of the X-ray gas over the volume of the radio lobes is 3.0$\\times10^{-12}$ Pa, a factor of 30 lower than the minimum internal pressure of the lobes (9.1$\\times10^{-11}$ Pa). This is not surprising given the small physical size of the radio galaxy. Independent evidence in support of a cluster comes from the gravitational arc detected with the HST (Deltorn et al. 1997) at 4.9 arcsec to the northeast of the radio galaxy (Fig. \\ref{fig:3c184hst}). These authors detect an over-density of galaxies around the radio galaxy 3C\\,184 and measure 11 galaxies at a redshift of $\\approx$1. From the measured redshift they derived a velocity dispersion of $\\sigma_v=634^{+206}_{-102}$ km s$^{-1}$. As a first approximation, we can estimate the expected temperature of the gas from the galaxy velocity (i.e. Sarazin (1986): $T_g \\simeq 7\\times10^7 {\\rm K} ~(\\sigma_v/1000$ km s$^{-1})^2$, which gives $T_g \\simeq 2.4^{+1.9}_{-0.7}$, in good agreement with our spectral results. The comparison with the independent analysis of Deltorn et al. (1997) confirm that the temperature should be around 3 keV, despite the large uncertainties due to the limited statistics in our spectral analysis. Deltorn et al.~ computed a mass enclosed within 40 $h_{70}^{-1}$ kpc (5 arcsec) of $\\sim2.1\\pm0.9\\times10^{13} h_{70}^{-1}$ $M_{\\odot}$ and derived a virial mass of 7.7 $\\times 10^{14}$ $h_{70}^{-1}$ $M_{\\odot}$, from the central velocity dispersion calculated for an isothermal sphere. If we adopt the proton density calculated in Sec. \\ref{sec:radprof1}, we estimate a gas-mass within a cylinder of radius 5 arcsec integrated along the line of sight of $\\sim1.2\\times 10^{11}$ $M_{\\odot}$. We compared this value with the total mass calculated in the same cylinder (and scaled for our cosmology) of Deltorn et al. (1997), e.g. $\\sim(2.1\\pm0.9)10^{13}$ M$_{\\odot}$, and found a gas-mass to total-mass ratio of order 0.01. This is at least a factor of 10 lower than that observed for clusters of galaxies at lower redshift. Errors on this estimate are large, and the largest uncertainty is likely to come from the $\\beta$-model parameters. The best estimate for the core-radius, seems to suggest a rather shallow potential well with respect to other clusters at the same redshift (e.g. Vikhlinin et al. 2002), more similar to galaxy groups, rather than rich clusters. On the other hand, the optical analysis of Deltron et al. supports the existence of a large total mass around 3C\\,184, and the temperature we measure also points in this direction. Since the spectral fit also implies an underlumious object, we can speculate that the cluster around 3C\\,184 is somehow peculiar in its relatively high temperature but low luminosity (and low gas fraction) if compared to lower redshift galaxy clusters. However, we note that Croston et al. (2003) found a trend for galaxy groups harbouring a radio-loud source to be hotter than radio-quiet galaxy groups. 3C\\,184 might show the effect of gas heating by a radio source at $z = 1$. The fact that the radio source is small and young and expanding might represent another piece of evidence in support to this picture. \\subsection{Component separation in 3C\\,292} At a redshift of $z=0.7$, 3C\\,292 is the nearest source in our small sample. For this reason we can better separate the X-ray components spatially. The spectral properties of the core of 3C 292 reveal an absorbed component which probably indicates a hidden quasar. With the pn camera, we also detect soft emission, and the normalisation of this component (0.13 nJy) supports the interpretation to be a soft radio-related emission from the core. The derived 0.2-10 keV flux (1.3$\\times10^{-14}$ \\fluxunit) is in agreement with the ROSAT-derived flux-flux correlation for radio/X-ray cores (Hardcastle \\& Worrall 1999). We have evidence for the presence of extended X-ray emission. Most of it is spatially correlated with the radio lobes, but we detected an additional extended component that we interpret as arising from a hot, cluster-like environment, and detect spatially out to a distance of 0.7 Mpc (100 arcsec) from the centre of the radio galaxy. Our best-fitting temperature, k$T$ = 2.2 keV, and bolometric luminosity, $L_{\\rm X}$ = 6.5$\\times10^{43}$ erg s$^{-1}$, place the cluster slightly above, but consistent with, the $L_{\\rm X} - T$ relation for clusters at $z>0.4$ (Vikhlinin et al. 2002) The spatial correlation between radio and X-ray emission, at the position of the radio lobes, strongly supports IC scattering as the physical process responsible for the X-ray emission aligned with the radio structure. The larger spatial extent of the radio emission can be explained by lifetime effects. However, some doubts arise on a pure non-thermal origin because a thermal model is a good fit to the data, and there is some evidence for line Fe L and Fe K emission. Let us assume that the X-ray lobes emit thermally and their temperature is 5.3 keV, to be compared to a 2 keV temperature of the gaseous environment as obtained from the spectrum extracted on larger scales and by excluding the lobe X-ray emission. A possible physical interpretation in this case is that the gas surrounding the radio lobes has been shocked by the expansion of the radio jet, as observed in the nearby source Cen A (Kraft et al. 2003). If this is the case, we are observing this phenomenon at high redshift for the first time. Although this effect is difficult to quantify, some contamination from the core via the wings of the PSF could contribute to artificially increasing the temperature, if a thermal model is adopted. On the other hand, the adoption of IC scattering as the physical model to explain the X-ray emission is also supported by requiring the radio plasma to be only within a factor of 2 of equipartition. Under the simple assumption that the cluster is described by an isothermal $\\beta$-model, we estimate that the pressure of the external medium at the distance of the radio lobes (2.5$\\times10^{-13}$ Pa) and the pressure in the radio lobes (7.9$\\times10^{-13}$ Pa) are of the same order, suggesting equilibrium of the radio-galaxy lobes with the external environment. \\subsection{3C 322} The poor statistics do not allow us to constrain well the origin of the X-ray emission. We have a detection, with each of the three cameras, and we observe a correspondence between the X-ray emission and the radio emission from the galaxy. There is an interesting correlation between the southern hotspot and the X-ray southern peak of the emission, while the northern hotspot is 10 arcsec~ (85 kpc) from the northern X-ray peak. However, the estimated SSC X-ray emission for hotspots is negligible with respect to the estimated IC emission from the lobes. Despite our reservations concerning the spectral fitting, the good agreement between the X-ray flux estimated at 1 keV by adopting a single power-law model with $\\Gamma$=1.6 and the expected IC flux from the lobes supports the IC interpretation. If we interpret the soft emission as thermal and due to a cluster environment, our spectral results would lead to a temperature (1.4 keV) and X-ray luminosity (1.5$\\times10^{44}$ erg s$^{-1}$), consistent with the $L_{\\rm X} - T $ relation when evolution is considered (Vikhlinin et al. 2002). Core emission from this source is not obviously detected, and a second component to fit the spectrum is not required. However, if the hard spectrum is fitted, analogously to 3C\\,292 and 3C\\,184 with an absorbed power law, we obtain similar absorption and photon index as for the other two sources." }, "0405/astro-ph0405269_arXiv.txt": { "abstract": "We explore isothermal shock formation in non-equatorial, adiabatic accretion flows onto a rotating black hole, with possible application to some active galactic nuclei (AGNs). The isothermal shock jump conditions as well as the regularity condition, previously developed for one-dimensional (1D) flows in the equatorial plane, are extended to two-dimensional (2D), non-equatorial flows, to explore possible geometrical effects. The basic hydrodynamic equations with these conditions are self-consistently solved in the context of general relativity to explore the formation of stable isothermal shocks. We find that strong shocks are formed in various locations above the equatorial plane, especially around a rapidly-rotating black hole with the prograde flows (rather than a Schwarzschild black hole). The retrograde flows are generally found to develop weaker shocks. The energy dissipation across the shock in the hot non-equatorial flows above the cooler accretion disk may offer an attractive illuminating source for the reprocessed features, such as the iron fluorescence lines, which are often observed in some AGNs. ", "introduction": "A realistic model of the AGN central engine may include the effects of the magnetic field. The black hole magnetosphere is studied first by \\cite{BZ77} in the context of winds and jets from radio-loud AGNs. The work has been extended to magnetospheric physics of accreting AGNs by various authors \\citep[e.g.,][hereafter TRFT02; Rilett et al.2004, in preparation]{Phinney83,TNTT90,TT01,TRFT02}. In this case, plasma particles should be {\\it frozen-in} to the magnetic field lines, and hence the accreting fluid should fall onto the black hole from regions {\\it above} the equatorial plane along the field lines\\citep[see, e.g., Figure 2 of TRFT02; Figure 1 of][]{TT01}. Then, magnetohydrodynamics (MHD) should become important to describe the motion of the particles associated with the background field. The resulting relativistic MHD shocks are explored by TRFT02 and Rilett et al. (2004, in preparation). In general, the MHD shocks can be hydro-dominated or magneto-dominated \\citep[hereafter T02]{Takahashi00,Takahashi02}. Obviously the MHD case is, however, very complicated. Therefore, the main motivation of our current paper is, as a starting point, to investigate the hydrodynamic limit, which should be valid in the case of small magnetization. Here, we adopt a model which should apply to the hydro-dominated shocks where the magnetic field does not make a significant contribution to the properties of the shocks under a weak field limit, because in such a case hydrodynamics should primarily control the shock formation. One-dimensional (1D), hot accretion flows around a black hole, generally treated as ideal hydrodynamical fluid, have been investigated by various authors \\citep{Sponholz94,Kato96,Kato98}. It has been found that such accretion flows must be transonic, and will become supersonic before reaching the event horizon, while it is subsonic at infinity. Once such a fluid becomes supersonic, it is likely that a standing shock wave will develop when shock conditions are met. There are roughly three types of standing shocks: adiabatic (Rankine-Hugoniot) shocks, isothermal shocks, and isentropic compression waves \\citep{Abramowicz90}. The first attempt was made by \\cite{Yang95} to self-consistently study the relativistic isothermal shock formation around black holes. In the case of isothermal shocks, the postshock fluid can lose substantial energy and entropy across the shock while the fluid temperature is continuous across the shock location \\citep{Lu97b}. \\citet[hereafter LY98]{Lu98} examined the isothermal shock formation in one-dimensional (1D) adiabatic hot flows in the Kerr geometry, for various flow parameters including black hole rotation. In our current paper, we extend the work by LY98 on isothermal shock formation in 1D adiabatic flows in the equatorial plane, to two-dimensional (2D) calculations for flows {\\it above} the equatorial plane, to investigate geometrical effects. One of our major motivations is to explore the possibility that shocks produced in such flows act as a high energy radiation source for some reprocessed features, such as the iron fluorescent lines, which are observed from some AGNs. It is generally considered that a hot illuminating source {\\it above} the cooler disk is required to produce such features \\citep{Fab89}. Section \\S{2} introduces our basic equations and assumptions. The hydrodynamic fluid equations are solved in the context of general relativity. We discuss the isothermal shock conditions and the stability of shocks. The results are presented in Section \\S{3} where we display shocks for various representative values of angular position, fluid energy, angular momentum, and black hole spin. Discussion and concluding remarks are given in the last section, \\S{4}. \\placefigure{fig:disk}% ", "conclusions": "Although our current work is partially motivated by the work of LY98, it also originates, in important ways, from our very recent work on black hole magnetospheres in accretion-powered AGNs - TRFT02 and Rilett et al. (2004, in preparation), where we explored adiabatic, relativistic MHD shocks produced in the black hole magnetosphere. We showed that strong shocks can indeed be formed in such flows for various relevant choices of flow parameters. In the current paper we extended these previous studies to explore the relativistic hydrodynamic flows which should apply to the case of weak magnetization. The reason is that the physics involved in the exact MHD case would be far more complicated, as noted already by TRFT02. Owing to these complexities, it was not straightforward for us to study the exact global shock-included MHD accretion flow solutions further in detail in a wider parameter/solution space. As our next step, in the current paper we assumed a simple model of conical accretion flows, due to the fact that modelling more realistic flows (such as the ones in hydrostatic equilibrium) would be very complicated, especially in the framework of the relativistic non-equatorial accretions considered here. An appropriate force balance under the general relativistic geometry should be taken into account in more sophisticated models, but that is beyond the scope of our present work. We, however, would like to stress that our results still represents (at least qualitatively) important physical characteristics of shock formation in non-equatorial accretion flows. Also, in the presence of the magnetosphere, the fluid particles would be frozen-in to the field and hence flow along the field lines. For such a situation, application of conventional thick accretion disk (or torus) models is not appropriate. We considered only the inner range of the shock formations in regions relatively close to the central engine, because of our interests in some X-ray observations of the reprocessed emission, such as the iron fluorescence lines, from some AGNs. Our results may offer the possibility of a high energy source in various broad regions ($0\\degr < \\theta_{sh} < 90\\degr$) above the disk plane. However, we find that the strongest shocks (high $M_1/M_2$ with large $E_{sh}$) should develop near the equator (large $\\theta_{sh}$), although the quasi-polar shock (small $\\theta_{sh}$) is also possible. We find no shocks in the polar region ($\\theta_{sh}=5\\degr$) when the preshock fluid energy is relatively large although the shock-free accretion is physically allowed. The magnitude of the energy release from the shock roughly increases as the shock location gets closer towards the black hole, as already found by LY98 by their 1D studies. In addition, however, we further find, from our 2D calculations, that the shock-induced energy release $E_{sh}$ greatly depends, not only on the fluid energy $E_1$ and angular momentum $\\ell$, but also on {\\it the angle of the shock location $\\theta_{sh}$}. The average shock strength (thus energy release) tends to be weaker towards the polar region. Although TRFT02 considered adiabatic shocks and hence no energy dissipation, in the current paper we adopted the isothermal shocks because that ensures a substantial amount of energy release at the shock locations. Also, through the stability analysis our energy source (i.e., the shock) is found to be stable. Although we adopted adiabatic flows in our current work, it may be noted that \\cite{DPM03} recently explored the isothermal shock formations in the isothermal flows, by adopting pseudo-Schwarzschild gravitational potentials and 1D equatorial flows. These authors concluded that their shocks also will release substantial amount of energy, which could be physically sufficient to become a radiation source for a strong X-ray flare. In the present work, it is argued that a single isothermal shock formation between the two sets of accreting flows (preshock/postshock flows) is very likely with a substantial amount of energy dissipation. One may also ask whether it is possible to have a sequence of shock formations one after another in the course of the accretion. For example, a preshock flow with energy $E_1$ gets shocked at a shock location $r_{sh,1}$ becoming a subsonic postshock flow with energy $E_2$. The same flow with energy $E_2$ then becomes supersonic and develops another shock at $r_{sh,2}$ where $r_{sh,2}